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UNIVERSITA’ DEGLI STUDI DI MILANO-BICOCCA Scuola di Dottorato di Scienze Corso di Dottorato di Ricerca in Fisica e Astronomia XVIII ciclo UNIVERSITÉ DE PROVENCE AIX-MARSEILLE I Ecole Doctorale ”Physique et Sciences de la Matière” Doctorat en Rayonnement et Plasmas A.A.2004-2005 ENVIRONMENTAL EFFECTS ON GALAXY EVOLUTION IN NEARBY CLUSTERS Coordinatore del Dottorato: Prof. Claudio Destri Directeur de l’École Doctorale: Prof. Jean-Jacques Aubert Tutore: Prof. Giuseppe Gavazzi Directeur de thèse: Dott. Alessandro Boselli Commissione-Jury: Dott. A. Boselli (Laboratoire d’Astrophysique de Marseille) Prof. V. Buat (Université de Provence) Prof. G. Gavazzi (Università di Milano - Bicocca) Prof. F. Haardt (Università dell’Insubria) Rapporteurs: Prof. C. Balkowski (Observatoire Astronomique Paris-Meudon) Dott. B. Poggianti (Osservatorio Astronomico di Padova) Tesi di Dottorato di: Luca Cortese Matricola R00280 ”Objectivity cannot be equated with mental blankness; rather, objectivity resides in recognizing your preferences and then subjecting them to especially harsh scrutiny ...and also in a willingness to revise or abandon your theories when the tests fail (as they usually do).” Stephen Jay Gould Acknowledgments This work represents the end point of my student career. After approximately twenty one years from my first entrance in a class room (it was September 1984 in Phoenix, AZ), I’m finally going to attend my last ”school” examination. Therefore I want to seize this opportunity in order to briefly remember and to thank some of the friends met during this journey. First of all I must thank my advisor Peppo Gavazzi, my scientific father, for his precious guidance and his teachings especially at the beginning of my research carrier. Special thanks to Alessandro Boselli, my co-advisor, first of all for the the last year spent in Marseille: a splendid experience. Thanks also for all his helpful advices, comments and supports on this and other works during the last three years. Many people contributed, directly or indirectly, to this work, and I am grateful to all of them. Merci beaucoup to Samuel Boissier for all the interesting discussions and, above all, for his precious lessons of French. Thanks to Veronique Buat for his help during the year spent in Marseille and for having initiated me in the obscure secrets of dust. Thanks to Barry Madore for his hospitality at the Carnegie Observatories, for his kindness, support and, especially, for his help in improving my written English. Muchas gracias to Armando Gil de Paz for his precious help on making the GALEX data available to me: without his contribution a great part of this work would not have been possible. Many thanks to Bianca Poggianti for a careful reading of my thesis and for her useful comments and suggestions. I would like also to thank Monica Colpi for her scientific and, especially, financial support during these three years. Arigato to Tsutomu Takeuchi and Akio Inoue for useful discussions about dust and galaxy evolution, for their kindness and help during my stay in Marseille and for having introduced me to Japanese cuisine. Many friends made the last three years unique. At Milano University life wouldn’t have been as much fun without all Peppo’s students. In particular thanks to Ilaria, Lea and Paolo for their unique support and thanks also to Chri for having installed Linux on my laptop, making me able to write this work. In Marseille thanks a lot to all the ”Café du Coin”: Helene, Claude, Kassem, Peter, Fabrice and the others. Thanks for all the coffees and cakes, and for having received me with open arms even if I wasn’t able to speak French. Thanks to Alexie, Jeanv vi Baptiste, Hector who shared the office with me, and a special thanks to Celine for having borne my never ending phone calls with my advisors, for her kindness and for her precious help in understanding french bureaucracy. Life in Marseille would have been completely different without the volley matches with Raph, Patrick, Seb, Mika, Fabrice and all the others. Finally, nothing of this would have been possible without the constant support of my parents and my brother Claudio, who have always encouraged me to continue this beautiful adventure. This research was partly supported the Università Italo-Francese through the Vinci Programme and by the CNES through GALEX-Marseille. Abstract The environmental effects on galaxy evolution in nearby clusters are investigated using a multiwavelength dataset. The present analysis is focused on the properties of three (Abell 1367, Virgo and Coma) among the best studied clusters in the local Universe. Due to the variety of their environmental conditions (e.g. spiral fraction, X-ray luminosity, evolutionary stage) they represent the most suitable ”laboratory” for comparative studies. By combining for the first time GALEX UV observations with optical, near and far infrared data, the evolutionary history of cluster galaxies is studied. The main goals of this thesis are: (a) The study of the dependence of the UV emission of galaxies from their morphological type, mass and the environment they inhabit, through the study of UV luminosity functions and color magnitude relations. (b) The study of UV dust extinction properties of local cluster galaxies and investigation of possible empirical relations useful to estimate the amount of UV attenuation in local and high redshift galaxies. (c) Investigation of the effect of large scale structures assembling on galaxy evolution through the dynamical analysis of Abell 1367, one of the best examples of a dynamically young local cluster of galaxies. (d) The characterization of the effects of different environmental mechanisms (i.e. gravitation interactions, ram pressure, preprocessing) on the evolutionary history of cluster galaxies in order to gain more insight on the origin of the morphology-density and star-formation-density relations. The observational evidences presented in this work suggest that: (I) Giant ellipticals are an old, homogeneous population showing no or little evolution at least in the past 8 Gyr; unlike dwarf ellipticals which still contain young stellar populations. (II) The importance of different environmental mechanisms has changed during the age of the Universe. Tidal interactions and preprocessing probably dominated the past Universe and shaped part of the morphology-density relation during the phase of cluster accretion of small groups. Ram pressure dominates in today clusters and is surely affecting the star formation history of galaxies but with less influence on their morphology. (III) The heterogeneous class of S0s galaxies, from bulge dominated to the disky S0s, is not the result of a single transformation mechanism: if ram pressure is able to produce disk dominated S0s, tidal interactions (and thus preprocessing) are required to account for bulge dominated S0s. (VI) Different observational evidences vii viii confirm the presence of a correlation between the mean age of stellar populations and galaxy mass (downsizing effect). In the framework of the hierarchical model of galaxy formation, the origin of the downsizing effect remains unsolved. This clear observational evidences represents one of today’s main challenge for models of galaxy evolution. Riassunto In questo lavoro vengono analizzati gli effetti dell’ambiente sull’evoluzione delle galassie, utilizzando una base di dati multi-lunghezza d’onda. In particolare tutta quest’analisi é focalizzata sullo studio di tre differenti ammassi di galassie dell’Universo Locale: Abell1367, Virgo, Coma. Questi tre ammassi sono tra i piú studiati nell’Universo locale e, date le loro differenti proprietá (e.g. frazione di galassie a spirale, luminositá X, stadio evolutivo), rappresentano dei laboratori ideali per quantificare l’influenza dell’ambiente sull’evoluzione delle galassie. Combinando per la prima volta osservazioni ultraviolette del satellite GALEX a dati ottici, in vicino e lontano infrarosso viene ricostruita l’evoluzione delle galassie d’ammasso. I principali obiettivi di questa tesi sono: (a) Studiare il legame tra le proprietá dell’emissione UV delle galassie, il loro tipo morfologico, la loro massa e l’ambiente in cui esse si trovano, attraverso l’analisi delle funzioni di luminositá UV e delle relazioni colore-magnitudine. (b) Comprendere le proprietá delle polveri interstellari responsabili dell’assorbimento della radiazione ultravioletta e ricavare relazioni empiriche utili per poter quantificare l’assorbimento della radiazione ultravioletta in assenza di osservazioni in lontano infrarosso. (c) Analizzare se e come lo stato dinamico di un ammasso é in grado di influenzare la storia evolutiva delle galassie, attraverso lo studio dell’ammasso di Abell1367: uno dei migliori esempi di ammasso locale, dinamicamente ancora giovane. (d) Quantificare l’influenza di diversi effetti d’ambiente (i.e. interazioni gravitazionali, ram-pressure, galaxy preprocessing) sull’evoluzione delle galassie d’ammasso, in modo da comprendere le origini del fenomeno di segregazione morfologica. Tutte le evidenze osservative presentate e analizzate in questo lavoro suggeriscono che: (I) Le ellittiche giganti rappresentano una popolazione vecchia, omogenea che non ha subito una significativa evoluzione negli ultimi 8 Gyr; al contrario dell’ellittiche nane che sono ancora oggi dominate da popolazioni stellari giovani. (II) L’influenza dell’ambiente sull’evoluzione delle galassie cambia sensibilmente con l’etá dell’Universo. Le interazioni gravitazionali ed il galaxy preprocessing sono stati gli effetti dominanti nell’Universo passato e sembrano essere i responsabili, almeno in parte, del fenomeno di segregazione morfologica. La ram pressure sembra essere dominante negli ammassi di oggi. Questo meccanismo é sicuramente in grado di influenzare la storia di forix x mazione stellare delle galassie, ma ha pochi effetti sulla loro morfologia. (III) Le galassie lenticolari (S0) risultano essere cosı́ il prodotto di processi completamente differenti: se oggi la ram pressure é in grado di trasformare una galassia a spirale in una lenticolare con piccolo bulge, sono necessarie interazioni gravitazionali per produrre i grandi bulge osservati in molte lenticolari nell’Universo locale. (IV) Diverse, e indipendenti, evidenze osservative confermano l’esistenza di una forte correlazione tra la massa degli oggetti e l’etá media della loro popolazione stellare (downsizing effect). Nel quadro dei modelli gerarchici di formazione delle strutture, l’origine del downsizing effect é tutt’ora sconosciuta. La comprensione di questo fenomeno rappresenta dunque una delle maggiori sfide per l’astronomia extragalattica. Résumé Ce travaille est dédié à l’étude des effets d’environnement sur l’évolution des galaxies dans l’Univers voisin, en utilisant un échantillon multi-longueur d’onde. En particulier toute cette analyse est focalisée sur les propriétés des trois différents amas des galaxies: Abell1367, Virgo et Coma. Ces trois amas des galaxies sont parmi les mieux étudiés dans l’Univers local et, en raison de la variété des leurs propriétés (par exemple fraction des galaxies à spirale, luminosité X, état dynamique), ils représentent des laboratoires, les plus appropriés, pour des études comparatives. En combinant pour la première fois des observations UV de GALEX à des données en optique, en voisin et en lointain infrarouge j’ai déterminé l’histoire evolutive des galaxies dans les amas. Les buts principales de cette thèse sont: (a) Étudier la variation des propriétés UV des galaxies en fonction des propriétés de l’environnement ou elles se trouvent, de leur masse et type morphologique, en analysant les fonctions de luminosité en UV et les relations couleur-magnitude. (b) L’étude du taux d’absorption des photons UV par les poussiéres interstellaires, pour obtenir des relations empiriques trés utils pour quantifier l’attenuation par poussiéres quand les données en infrarouge lointain sont absentées. (c) Analyser l’effet de la formation des amas sur l’évolution des galaxies en étudiant l’amas d’Abell1367, un des meilleurs exemples d’amas voisin et dynamiquement jeune. (d) Comprendre l’influence des différents effets d’environnement sur l’histoire evolutive des galaxies d’amas, pour comprendre l’origine de la ségrégation morphologique dans les amas. Touts les résultats obtenus dans ce travaille montrent que: (I) La population des galaxies elliptiques géants est vieille et homogène. Elle ne montre pas d’évolution au moins dans les derniéres 8 Gyr; au contraire des elliptiques naines qui contiennent toujours populations stellaires jeunes. (II) L’importance relative des différents mécanismes d’environment varie avec l’âge de l’Univers. Les interactions de marée et le prepocessing ont probablement dominées dans l’Univers passé et ont contribuées (au moins en partie) à la ségrégation morphologique, pendant la formation des amas par des petits groupes des galaxies. La pression dynamique domine dans les amas d’aujourd’hui et elle affecte sûrement l’histoire de formation des etoiles des galaxies avec moins d’influence sur leur morphologie. (III) La classe hétérogène des galaxies S0s (lenticuliers), n’est pas le résultat d’un seul mécanisme de transformation: si la xi xii pression dynamique peut produire S0s dominées par le disque, les interactions de marée (et le preprocessing) sont exigées pour expliquer les S0s dominées par le bulbe. (IV) Différentes évidences suggèrent la présence d’une corrélation entre l’âge moyen des populations stellaires et la masse des galaxies (downsizing effect). Dans le cadre du modèle hiérarchique de formation des galaxies, l’origine de cet effet n’est pas encore résolue. Il représente aujourd’hui une des défis pour les modèles d’évolution des galaxies. Contents 1 Introduction 5 2 GALEX & GOLDMINE: A multiwavelength window on the Universe 2.1 The Galaxy Evolution Explorer . . . . . . . . . . . . . . . . . . 2.1.1 The Prime Mission . . . . . . . . . . . . . . . . . . . . . 2.1.2 Data collection mode . . . . . . . . . . . . . . . . . . . . 2.1.3 Counts vs. magnitudes and fluxes conversions . . . . . . 2.2 The Galaxy On Line Database Milano Network . . . . . . . . . 3 The 3.1 3.2 3.3 3.4 Local . . . . . FAUST-FOCA UV luminosity function of nearby clusters Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . The Data . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . The UV luminosity functions . . . . . . . . . . . . . . . . . . . . 3.3.1 The composite cluster luminosity function . . . . . . . . . Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 4 GALEX UV luminosity function 4.1 Introduction . . . . . . . . . . . 4.2 UV data . . . . . . . . . . . . . 4.3 The luminosity function . . . . 4.4 Discussion . . . . . . . . . . . . of Abell1367 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 5 Multiple merging in Abell1367 5.1 Introduction . . . . . . . . . . . . . 5.2 Observations and data reduction . . 5.3 The global velocity distribution . . 5.4 Localized velocity structures . . . . 5.5 The cluster dynamics . . . . . . . . 5.5.1 The North-West subcluster . 5.5.2 The South-East subcluster . 1 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 15 15 16 17 18 20 . . . . . 23 23 24 25 26 28 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 33 33 33 37 40 . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 43 43 44 46 48 51 53 56 2 CONTENTS 5.6 5.7 5.8 5.9 Star formation activity in the infalling groups Cluster mass . . . . . . . . . . . . . . . . . . . Two-Body Analysis . . . . . . . . . . . . . . . Conclusions . . . . . . . . . . . . . . . . . . . 6 Unveiling the evolution of early type galaxies 6.1 Introduction . . . . . . . . . . . . . . . . . . . 6.2 Data . . . . . . . . . . . . . . . . . . . . . . . 6.3 The UV properties of early-type galaxies . . . 6.4 Discussion and conclusion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . with GALEX. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 7 UV dust attenuation in normal star forming galaxies 7.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . 7.2 The Data . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 7.2.1 The optically-selected sample . . . . . . . . . . . . . . 7.2.2 The starburst sample . . . . . . . . . . . . . . . . . . . 7.3 The LT IR /LF U V − β relation for normal star-forming galaxies 7.3.1 The dependence on the birthrate parameter . . . . . . 7.4 A(Hα) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 7.4.1 Estimate of A(Hα) . . . . . . . . . . . . . . . . . . . . 7.4.2 The β-A(Hα) relation . . . . . . . . . . . . . . . . . . 7.5 Relations between dust attenuation and global properties. . . 7.5.1 Metallicity . . . . . . . . . . . . . . . . . . . . . . . . . 7.5.2 Luminosity . . . . . . . . . . . . . . . . . . . . . . . . 7.5.3 Surface brightness . . . . . . . . . . . . . . . . . . . . . 7.5.4 LHα /LF U V ratio . . . . . . . . . . . . . . . . . . . . . . 7.6 A cookbook for determining LT IR /LF U V ratio . . . . . . . . . 8 High velocity interaction: NGC4438 in the Virgo cluster 8.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . 8.2 Data . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 8.3 The UV emission and the star formation history of NGC 4438 8.4 Discussion and conclusion . . . . . . . . . . . . . . . . . . . . 9 Ram Pressure stripping: NGC4569 in the Virgo cluster 9.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . 9.2 Data and models . . . . . . . . . . . . . . . . . . . . . . . . 9.3 The star formation history of NGC 4569: model predictions 9.4 Discussion and conclusion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 59 61 62 65 . . . . 69 69 70 71 75 . . . . . . . . . . . . . . . 81 81 84 84 86 86 89 90 90 92 94 94 97 99 101 103 . . . . 107 107 108 110 113 . . . . 117 117 118 120 121 CONTENTS 10 Galaxy Pre-processing: the blue group infalling in Abell1367 10.1 Introduction . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.2 Observations . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.2.1 HI observations . . . . . . . . . . . . . . . . . . . . . . . . . . 10.2.2 UV to near-IR imaging . . . . . . . . . . . . . . . . . . . . . . 10.2.3 Hα imaging . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.2.4 MOS spectroscopy . . . . . . . . . . . . . . . . . . . . . . . . 10.2.5 High Resolution spectroscopy . . . . . . . . . . . . . . . . . . 10.3 Results . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.3.1 Kinematics . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.3.2 Hα properties . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.3.3 HI properties . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.3.4 The fate of the stripped gas . . . . . . . . . . . . . . . . . . . 10.3.5 The metal content . . . . . . . . . . . . . . . . . . . . . . . . 10.3.6 Dating the starburst. . . . . . . . . . . . . . . . . . . . . . . . 10.3.7 CGCG97-120: simply a foreground galaxy, or an high velocity intruder? . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.4 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.4.1 The evolutionary history of the Blue Infalling Group . . . . . 10.4.2 The contribution of preprocessing to cluster galaxies evolution. 3 127 127 128 128 131 131 131 134 135 135 138 143 148 149 151 156 157 157 158 11 Discussion & Conclusions 165 11.1 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 165 11.2 Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 175 A The extinction correction 179 B Estimate of the < 912Å flux from Hα + [NII] 181 Bibliography 183 Chapter 1 Introduction Eighty-five years are as short as a jiffy compared to the whole history of humanity and science, but this is the brief time men needed to upset their view of the Universe they inhabit. Let us return for a moment at the beginning of this story: July 26, 1920, Harlow Shapley and Herber Curtis confront their positions on the size of the Universe and the nature of the spiral nebulae in talks later called the Great Debate (see Trimble 1995, for a review). Curtis argued that the Universe is composed of many galaxies like our own, which had been identified by astronomers of his time as spiral nebulae. Shapley argued that these spiral nebulae were just nearby gas clouds, and that the Universe was composed of only one big Galaxy: our Milky Way. The resolution of the debate came in the mid 1920’s. Using the 100 inch telescope at Mount Wilson, Edwin Hubble identified Cepheid variable stars in the Andromeda Galaxy (M31). These stars resulted far beyond the most distant stars known in our galaxy and allowed Hubble (1925) to show that M31 was a galaxy much like our own. With this discovery, the known universe expanded immensely and, in the same time, a new research area was born: extragalactic astronomy. Thanks to overwhelming technological progress, during its first ∼85 years of life, extragalactic astronomy has provided us with a detailed description of the Universe from our neighbours (the Local Group) to its observable edges (the Cosmic Microwave Background). We know that most of the visible matter in the Universe, in the form of stars, gas, and dust grains, is organized in galaxies. Galaxies come in many different forms and sizes (as clearly shown in Fig.1.1), but they can be broadly divided into two main species. Spirals, with a flattened, disk-like shape, blue colors, much gas and dust, and a widespread star formation activity that results in the presence within them of many young stars. Ellipticals, with a spheroidal shape, red colors, little or no gas and dust, and no star formation activity, thus containing exclusively old stars. We also know that the density of galaxies in the local Universe is not at all constant, but it spans from ∼ 0.2 ρ0 in voids to ∼ 5 ρ0 in superclusters and filaments, ∼ 100 ρ0 in the cores of rich clusters, up to ∼ 1000 ρ0 in compact groups, where ρ0 is the average 5 6 1. Introduction Figure 1.1: An example of the heterogeneous population of galaxies that inhabit our Universe. Mosaic of RGB (g,r,i) images adapted from Frei et al. (1996) 7 “field” density (Geller & Huchra 1989). It is well established that morphological type and local density are not independent quantities. In their analysis of 55 nearby clusters, Dressler (1980) and Whitmore et al. (1993) demonstrated that the fraction of spiral galaxies decreases from 60% in the “field” to virtually zero in the cores of rich clusters, compensated by an opposite increase of elliptical and S0 galaxies. This phenomenon, known as morphology segregation, is considered as the clearest observational signature of significant environmental dependences of the processes that govern the formation and the evolution of galaxies. Understanding the origin of this phenomenon (”Nature or Nurture?”) probably represents one of the major challenges of extragalactic astronomy. One possible way to overcome this problem is to take advantage of the effect provided by the finite speed of light. Observing today galaxies at different distances means observing them at different epochs in the history of the Universe, and thus with different ages. This investigative method is providing us with a sort of evolutionary sequence for galaxies: starting from the pioneering work by Butcher & Oemler (1978, 1984) we know that distant (and thus young) cluster of galaxies contain a much higher fraction of blue galaxies than nearby clusters. Recently Dressler et al. (1997) used high-resolution imaging with the Hubble Space Telescope (HST) to measure the morphology-density relation in the core regions of a sample of rich clusters at z ∼0.5. They found that the fraction of lenticular galaxies (S0s) in clusters declined by a factor of 2-3 between z = 0 and z = 0.5, and this evolution was accompanied by a corresponding increase in the fraction of star-forming spirals (see also Couch et al. 1998; Treu et al. 2003). Many research groups have suggested that the predominance of early type galaxies in local clusters is the result of physical processes that suppress star formation and eventually alter galaxy morphology. and several mechanisms have been proposed (see Boselli & Gavazzi 2006, for a detailed review): • Galaxy interaction with the intra-cluster medium (ICM). Ram pressure stripping (Gunn & Gott 1972). As a galaxy orbits through a cluster, it experiences ram pressure from the ICM. When the ram pressure is greater than the binding force, the cold gas will be stripped (Abadi et al. 1999; Quilis et al. 2000; Vollmer et al. 2001). Before leading to complete gas ablation, ram-pressure could produces significant compression ahead of the galaxy temporally increasing its star formation activity (Bekki & Couch 2003). Even if it is well established that this phenomenon would finally lead to a gradual decrease in galaxy star formation activity, its effects on galaxy morphology are not yet completely understood (Fujita & Nagashima 1999; Mihos 2004a). Rampressure stripping is likely to be effective in the central region of clusters where the density of intra-cluster medium (ICM) is high. Viscous stripping (Nulsen 1982). In a galaxy travelling into the ICM the outer layers of the interstellar medium (ISM) experience a viscosity momentum trans- 8 1. Introduction fer that could be sufficient for dragging out part of its gas. Thermal evaporation (Cowie & Songaila 1977). If the ICM temperature is high compared to the galaxy velocity dispersion, at the interface between the hot ICM and the cold ISM the temperature of the ISM rises rapidly, thus the gas evaporate and is not retained by the galaxy gravitational field. Starvation (or strangulation) (Larson et al. 1980a). This mechanism consists in the removal of the diffuse hot gas reservoir that is confined in the galaxy halo. Since this tenuous halo is less bound than the cold gas in the disk, its stripping is considerably easier (Bekki et al. 2002). A galaxy whose hot gas reservoir is removed slowly, exhausts its cold gas in more than one gigayear, because there is no supply of fresh gas from the surrounding hot gas. Note that while stripping gas from disks induces a truncation of star formation activity on a short timescale (∼ 107 yr), strangulation is expected to affect a galaxy star formation history on a long time scale (> 1 Gyr) provoking a slowly declining activity which consumes the disk gas after the supply of cooling gas has been removed. All of the above mechanisms but starvation need relatively high density of hot intra-cluster gas, and thus likely to happen in the central region of clusters. However Fujita (2004) has pointed out that ram pressure and thermal evaporation could not be negligible in cluster sub-clump regions (small groups around a cluster). • Galaxy-galaxy gravitational interaction. Collisions or close encounters between galaxies can have a strong effect on their morphology and star formation rates. Various simulations have shown that major mergers between disk galaxies can produce galaxies resembling ellipticals as merger remnants (e.g.,Toomre & Toomre 1972; Barnes & Hernquist 1996) and that accretion of small satellites onto spirals can transform the host spiral to S0 type (Walker et al. 1996). The tidal forces generated during the interaction tend to funnel gas toward the galaxy center. It is likely that this will fuel a central starburst, ejecting a large fraction of material. Gas in the outer part of the disk, on the other hand, will be drawn out of the galaxy by the encounter (Mihos 2004a). Although individual collisions are expected to be most effective in groups because the velocity of the encounters is too high for such mergers to be frequent (Ghigna et al. 1998; Okamoto & Habe 1999), Moore et al. (1996) showed that the cumulative effect of many weak high velocity interactions (i.e. galaxy harassment) can also be important in cluster of galaxies. However its influence is largely limited to low luminosity galaxies, while in bright spirals its effects are considerably milder (Mihos 2004a; Moore et al. 1996). • Galaxy-Cluster gravitational interaction. Tidal compression of galactic gas via interaction with the whole cluster potential can effectively perturb cluster galax- 9 ies, inducing gas inflow, bar formation, nuclear and perhaps disk star formation (Merritt 1984; Miller 1986; Byrd & Valtonen 1990). On the other hand, gas can be hardly removed directly by the interaction (Boselli & Gavazzi 2006). Although we have collected a plethora of observational evidences that at least some of these processes are playing a significant role on galaxy evolution we have not shed light on the origin of the morphology density relation. This is in part due to the fact that we do not yet know their detailed physics and the relative importance of each mechanism during the different phases of galaxy evolution. Moreover the arduous effort of reconstructing the evolutionary history of galaxies would turn out to be completely useless if we did not take into account that the whole Universe is evolving, changing the physical condition of the environments populated by galaxies. In fact different and mostly independent observational evidences, as the Cosmic Microwave Background radiation (Kogut et al. 2003), the large scale structure (Hawkins et al. 2003) and supernovae observations (Tonry et al. 2003), are telling us that the Universe in not only expanding (Hubble & Humason 1931), but it is also accelerating. If theorists are right, this implies that the Universe is dominated by its energetic and a matter dark components, whose nature is still completely unknown. The dark energy term (usually indicated with the cosmological constant Λ) allows for the current accelerating expansion of the universe. Currently, ∼70% of the energy density of the Universe is supposed to be in this form. The dark matter component of the Universe is supposed to be cold (i.e. not thermalized), non-baryonic, collisionless ”material”. This component makes up ∼26% of the energy density of the present Universe and only the remaining ∼4% is the matter and energy we directly observe. The only way to shade light on the properties of our, mostly obscure, Universe is thus through numerical simulations (e.g. Kauffmann et al. 1993; Springel et al. 2005). In particular hierarchical galaxy formation (White & Rees 1978) models within a Λ cold dark matter (ΛCDM) cosmogony are currently considered the most successful paradigm for understanding the evolution of matter in the Universe. In this scenario, structures grow hierarchically via gravitational instability from small perturbations seeded in the early epoch. The density of dark matter its component is a proxy for the epoch of initial collapse of a given structure: the most massive structures at any epoch represent the earliest that collapsed (Springel & Hernquist 2003). After their collapse, structures grow up through infall of smaller groups (Kauffmann 1995). However the typical size of the infalling groups increases with the age of the Universe but their infall rate considerably decreases (Ghigna et al. 1998; Okamoto & Habe 1999; Gnedin 2003). This means that clusters have accreted great part of their galaxy population in the past, through infalling of small groups. Today the accretion of new members is supposed to be rare and to happen mainly through the merging of big subclusters. Adding the well known observational evidence that star formation rapidly decreases 10 1. Introduction with the age of the Universe (Lilly et al. 1996; Madau et al. 1998), we are facing a scenario that, at a first look, seems to suggest that studying star formation in rich clusters today is a melancholy affair. The Universe we inhabit today is old, and most of its star formation activity has gone out. In addition (and this is the worse part of the story) the Universe dramatically evolved itself, altering continuously the physical conditions of the environments populated by galaxies. This implies that galaxies could have experienced different environmental effects during their history and that the dominant process in the local Universe could have been completely negligible in the early stages of its evolution, while the process shaping galaxy evolution could be less important in today clusters. Let us imagine, as predicted by models, that a great fraction of today cluster galaxies have infalled, within a compact group, into a cluster ∼5 Gyr ago. While in the group environment tidal interactions were very strong and influenced significantly star formation activity and galaxy morphology; today, in cluster environment, gravitational interactions are less probable due the large relative velocities of cluster members (Ostriker 1980). Thus great part of galaxy evolution took place before the infalling into the cluster core. The discourage felt by a young student facing this music increases reading the recent review by Dressler (2004) on ”Star forming galaxies in clusters”: What we see in clusters today is only a faint echo of what once was... looking for star formation in today’s clusters is a little bit like searching for the last cashew in a picked-over nut-cup.[...] Star formation in rich clusters today is a pretty sad affair. Spirals are ”running down compared to half-a-Hubble time ago. The spirals that will be drawn into rich clusters in the future will die the death of a thousand cuts: in the rich group environment into which they have for so long been entrained, they are likely already to have had their fates sealed long ago. Thus, what has he to do? Give up and concentrate all his efforts on the study of the high redshift, still young, Universe? Obviously the answer is no; and not because this work would be useless. High redshift and local observations are complementary to give more insights on galaxy evolution and, until we will be able to understand all the physical mechanism influencing the present evolution of nearby galaxies (and we are still far from reaching this goal), it would be an error to concentrate all our efforts only at high redshift. Observations of the high redshift Universe approach us to the mechanisms that maybe shaped part of the morphology density relation; however today there is still insufficient high-quality data to put strong constraint on different models (Dressler 2004). On the contrary in the local Universe, maybe we are missing most of the action, but we have the unique possibility to observe in detail galaxy properties over the whole range of sizes and masses, and study in detail the effects of different environmental mechanisms. In particular, what makes the local Universe still exciting? What can we learn about galaxy evolution that would still be impossible if we moved to higher redshifts? Owing 11 to the high quality images we can obtain for local galaxies, an extremely accurate and homogeneous morphological classification is possible down to MB ≤ −13, allowing a detailed discrimination among different subclasses of early-type galaxies (ellipticals, lenticulars, dwarfs) and among early-type galaxies and quiescent spirals (see the Virgo Cluster catalogue a sort of ”milestone” of the morphological classification, Binggeli et al. 1985). Accurate morphological classification becomes a difficult task just at the Coma cluster distance (z ∼0.025) and more or less impossible at higher redshift (Abraham et al. 1996a). The objects in the images are very small, thus it is very hard to detect the fine structure elements needed to distinguish different classes. In order to solve this problem alternative classifications based on structural parameters (Abraham et al. 1996b) or on spectral type (Madgwick et al. 2002) have been proposed but they are only useful to discriminate between a star forming disk and a quiescent bulge dominated galaxy, completely failing to distinguish between an elliptical and a lenticular or between an early-type galaxy and a bulge-dominated Sa spiral disks (Scodeggio et al. 2002; Gavazzi et al. 2002a). Thus at high redshift we can observe the evolution of the star formation-density relation (the Butcher-Oemler effect) but we cannot investigate morphological transformations that eventually affected galaxy evolution (Smail et al. 1997; Fabricant et al. 2000; Smith et al. 2005). Moreover in the local Universe we can study galaxies spanning all ranges of mass and luminosity, reaching very faint (MB ∼ −13) low surface brightness (∼30 mag arcsec2 ) dwarf galaxies. This is crucial to study the (strong) dependence of galaxy evolution (Gavazzi et al. 2002a) and environmental effects with mass since the anti-hierarchical relation between star formation history and galaxy mass is one of the great challenge for models of galaxy evolution. Moreover dwarf galaxies today represent probably the major failure of hierarchical galaxy formation models: cold dark matter theory predicts that the groups and clusters of galaxies should contain many more dwarf objects than the observed number of dwarf galaxies (Klypin et al. 1999; Moore et al. 1999). Several explanation has been proposed (Somerville 2002), and even if no solution has been found so far, it is indisputable that the only way to solve this problem is to understand the formation and evolution of dwarf galaxies, a task possible only in the local Universe. Another serious limit of high redshift observations is the quantification of star formation activity in galaxies. The easiest and common way to estimate star formation rate for distant galaxies is through rest-frame ultraviolet (UV) observations. However ultraviolet emission is strongly affected by dust attenuation: absorption by dust grains reddens the spectra at short wavelengths and modifies altogether the spectral energy distribution of galaxies. Since the UV radiation is emitted by young stars (t < 108 yr) that are deeply embedded in dust clouds than older stellar populations, rest-frame UV observations can lead to incomplete and/or biased reconstructions of the star formation activity and star formation history of galaxies. Moreover we have not yet a good characterization of the dust attenuation properties in galaxies and of their 12 1. Introduction dependences with galaxy type (i.e. normal star forming galaxies vs. starburst) and no proper corrections have been achieved, having no possibility to correctly quantify the star formation rate at high redshift. As extensively discussed in Chapter 7 of this work, understanding dust properties and looking for empirical relations suitable for deriving dust attenuation corrections is today possible only for low redshift galaxies: in this case the study of the local Universe is mandatory to correctly interpret what we observe in distant galaxies. Finally, as remarked by Poggianti (2004a), in order to understand what happens to galaxies in clusters, two crucial pieces of information are 1) the gas content of cluster galaxies (i.e. the fuel for future star formation) and 2) the spatial distribution of the gas and of the star formation activity within each galaxy (i.e. differences from field galaxies are good indicators of environmental effects); and both can be achieved only in the local Universe. Neutral hydrogen (HI) and Hα observations observations 1 are still a prerogative of nearby galaxies. In the near future, thanks to the advent of the Arecibo L-band Feed Array, it should be possible to detect an hydrogen mass of ∼ 109 M at z ∼0.15, but only with a very high integration time (∼ 70 hours per beam). The few examples shown above represent only the tip of the iceberg of the unique capability of local Universe observations to disclose the secret of galaxy evolution. We would lose too much, without any significant improvement, if we abandoned observations of nearby galaxies in order to move our attention at high redshift galaxies. The aim of this work Firmly convinced of the great significance of nearby Universe observations, I have concentrated all my PhD work on the study of environmental effects on the evolution of nearby clusters. In particular this thesis will focus on three different clusters: Abell1367, Virgo and Coma. These three clusters are among the best studied in the local universe and, due to the variety of their environmental conditions (e.g. spiral fraction, X-ray luminosity, evolutionary stage) they represent the most suitable ”laboratory” for comparative studies. The novelty of this work is that in addition to the optical and near infrared observations carried out during the last fifteen years by G.Gavazzi and A.Boselli (available through the GOLDMine database, Gavazzi et al. 2003a: http://goldmine.mib.infn.it) I will take for the first time advantage of recent UV observations by the Galaxy Evolution Explorer (GALEX, Martin et al. 2005). The use of a multiwavelength dataset is crucial to understand galaxy evolution since different galactic components such as old, new or evolved stars; active galactic nuclei; the interstellar medium contribute in different amounts to the observed emission at different wavelengths, from the radio to X-rays. Therefore, the comparison of global emission properties at a wide range of wavelengths can give 1 The Hα Balmer emission (λ=6562.8 Å) is the most direct indicator of the current (< 4 106 yrs), massive (> 8 M) star formation activity in galaxies (Kennicutt 1998) 13 us precious insight on the relative importance of these components, as well as on the origin of some parts of the emission spectrum. Since different emission bands have different sensitivities to absorption, their comparison may also give us insight into the dust content of the emitting regions. Moreover, comparison of global multiwavelength emission properties of galaxies of different morphology can give us insight on the relative presence of different galactic components throughout the Hubble sequence. While most of the studies of galaxies make use of individual energy bands, mainly the optical but also the radio and, more recently, the X-ray and infrared, it is rarer to find work comparing data from two or more emission windows. In particular the rest frame UV emission provides a powerful tool for measuring and understanding star formation in galaxies at all epochs. Ironically, the interpretation of high-redshift galaxies in the rest UV is most limited by the lack of large, systematic surveys of lowredshift UV galaxies serving as a benchmark. However, before the launch of GALEX, only a few experiments had observed the nearby Universe at ultraviolet wavelengths (Smith & Cornett 1982; Lampton et al. 1990; Kodaira et al. 1990). Among them, the FOCA experiment (Milliard et al. 1991) allowed the first determinations of the UV LF of local field galaxies (Treyer et al. 1998; Sullivan et al. 2000) and and the local rest-UV anchor point for the star formation history plot. However its low sensitivity and resolution and the small sky area covered. With its large field of view (diameter ∼1.2 degrees), high sensitivity and two ultraviolet filters, GALEX has opened a new era in the UV astronomy, providing us for the first time with a large, complete and homogeneous dataset to study star formation activity in galaxies. Using this unique mine of data I will investigate the properties of galaxies from different points of view: one statistical, analyzing the global properties of the whole cluster sample; and another much more focused on the study of particular objects considered as prototypes of the different ways in which the environment could influence galaxy evolution. The comparison of all the observational results with models will be used to build up an evolutionary scenario for galaxies, linking the information I obtain in this work to what we know (or think to know) about the evolution of galaxies at higher redshifts. The organization of the thesis In Chapter 2 I briefly describe the different datasets used in this work: the GALEX satellite and its mission, and the GOLDMine database. In Chapter 3 and 4 I start the statistical analysis of cluster galaxies, computing the UV luminosity function for nearby clusters. The analysis presented in Chapter 3 was performed before the launch of GALEX, thus I used data from the FOCA (Milliard et al. 1991) experiment of the three nearby clusters studied in this work. When GALEX was launched I had the possibility to extend my analysis two magnitudes deeper with higher quality data. First of all, this double estimate allow me to directly 14 1. Introduction compare two different and independent datasets. Then the comparison between the cluster luminosity function and the field one is used to determine whether the environment affects the shape of the cluster luminosity function. In Chapter 5 I study the influence of the dynamical state of a cluster to the evolution of galaxies, performing a detailed dynamical study of the Abell cluster 1367. This cluster is considered as the prototype of a dynamically young local cluster, thus representing a good place to study the effects of a cluster’s assembly on galaxy evolution. Although X-ray, radio and optical observations suggest that Abell 1367 is dynamically young and it is still undergoing the process of formation, detailed spatial and dynamical analysis of this cluster has not been attempted so far. Since the dynamical state of a cluster is directly linked with its evolution this work will allow us to have a clear picture of the past, current and future assembly history of this structure and its galaxies. In Chapter 6 I focus my attention on the population of early-type galaxies in clusters, in particular studying the UV properties of giant and dwarf ellipticals and lenticulars in the Virgo cluster, in order to determine whether these different morphological types had the same evolutionary history or not. On the contrary from Chapter 7 till the end of this work I move my attention to the star forming cluster population. As discussed above if we want to use ultraviolet radiation to correctly estimate star formation we need to correct for dust attenuation. Thus in Chapter 7 I present an analysis of dust attenuation properties in nearby cluster star forming galaxies, obtaining a cookbook in order to estimate dust attenuation without far infrared observations. This analysis represents the tip of the iceberg and only a future comparison with different dust models will allow us to understand dust attenuation and to know how to correct UV observations of local and high redshifts galaxies. Thus, a statistical analysis of star formation activity in cluster galaxies using UV data is still impossible. For this reason in Chapter 8, 9 and 10 I will focalize my attention on the study of three particular cluster galaxies considered as the prototypes of the three main environmental effects observed in clusters: tidal interaction, ram pressure stripping and preprocessing, respectively. These unique astrophysical laboratories will help me to understand the effects of different physical mechanisms on galaxy evolution in more depth. Finally in Chapter 11 I will summarize the evolutionary scenario for cluster galaxies which emerged from this work. Great part of this thesis is published or submitted for publication on major astronomical refereed journals: Gavazzi et al. (2003b, 2006); Cortese et al. (2003a, 2004, 2005, 2006); Boselli et al. (2005a,b). Chapter 2 GALEX & GOLDMINE: A multiwavelength window on the Local Universe 2.1 The Galaxy Evolution Explorer The Galaxy Evolution Explorer (GALEX) is a NASA Small Explorer class mission. It consists of a 50 cm-diameter, modified Ritchey-Chrétien telescope with four operating modes: Far-UV (FUV) and Near-UV (NUV) imaging, and FUV and NUV spectroscopy. The telescope has a 3-m focal and the field of view is 1.2◦ circular (see Fig 2.1 and Table 2.1). Spectroscopic observations are obtained at multiple grism-sky dispersion angles, so as to mitigate spectral overlap effects. The FUV (1528Å: 13441786Å) and NUV (2271Å: 1771-2831Å) imagers (see Fig.2.2) can be operated one at a time or simultaneously using a dichroic beam splitter. The FUV detector is preceded by a blue-edge filter that blocks the night-side airglow lines of [OI]1304, 1356, and Lyα. The NUV detector is preceded by a red blocking filter/fold mirror, which reduces both zodiacal light background and optical contamination. The peak quantum efficiency of the detector is 12% (FUV) and 8% (NUV). The detectors are linear up to a local (stellar) count-rate of 100 (FUV), 400 (NUV) cps, which corresponds to mAB ∼ 14 − 15. The resolution of the system is typically 4.5/6.0 (FUV/NUV) arcseconds (FWHM), and varies by ∼20% over the field of view. Further detail about the mission, in general, and the performance of the satellite, in specific, can be found in Martin et al. (2005) and Morrissey et al. (2005), respectively. The mission is nominally expected to last 38 months; GALEX was launched into a 700 km, 29◦ inclination, circular orbit on 28 April 2003. 15 16 2. GALEX & GOLDMINE: A multiwavelength window on the Local Universe Figure 2.1: Cross section of the instrument portion of GALEX. The optical path is outlined in blue. Overall dimensions of the view shown are 1.5 m 1 m (adapted from Morrissey et al. 2005). 2.1.1 The Prime Mission GALEX is currently undertaking the first space UV sky-survey, including both imaging and grism surveys. The prime mission includes an all-sky imaging survey (AIS: 75-95% of the observable sky, subject to bright-star and diffuse Galactic background light limits) (mAB ' 20.5), a medium imaging survey (MIS) of 1000 deg2 (mAB ' 23), a deep imaging survey (DIS) of 100 square degrees (mAB ' 25), and a nearby galaxy survey (NGS). Spectroscopic (slit-less) grism surveys (R=100-200) are also being undertaken with various depths and sky coverage. Many of the GALEX fields overlap existing and/or planned ground–based and space-based surveys being undertaken in other bands. All-sky Imaging Survey (AIS): The goal of the AIS is to survey the entire sky subject to a sensitivity of mAB ' 20.5, comparable to the POSS II (mAB =21 mag) and SDSS spectroscopic (mAB =17.6 mag) limits. Several hundreds to 1,000 objects are in each 1 deg2 field. The AIS is performed in roughly ten 100-second pointed exposures per eclipse (∼10 deg2 per eclipse). Medium Imaging Survey (MIS): The MIS covers 1000 deg2 , with extensive overlap of the Sloan Digital Sky Survey. MIS exposures are a single eclipse, typically 1500 seconds, with sensitivity mAB ' 23, net several thousand objects, and are well-matched to SDSS photometric limits. 2.1. The Galaxy Evolution Explorer Item Bandwidth: Effective wavelength (λeff ): Field of view: Zero point (m0 ): Image resolution (FWHM): Spectral resolution (λ/∆λ): Detector background (typical): Total: Diffuse: Hotspots: Sky background (typical): Limiting magnitude (5σ): AIS (100 sec): MIS (1500 sec): DIS (30000 sec): 17 FUV Band 1344 – 1786 Å 1528 Å 1.28◦ 18.82 mag 4.5 arcsec 200 NUV Band 1771 – 2831 Å 2271 Å 1.24◦ 20.08 mag 6.0 arcsec 90 78 cnt sec−1 0.66 cnt sec−1 -cm−2 47 cnt sec−1 2000 c-sec−1 193 cnt sec−1 1.82 cnt sec−1 -cm−2 107 cnt sec−1 20000 cnt sec−1 19.9 mag 22.6 mag 24.8 mag 20.8 mag 22.7 mag 24.4 mag Table 2.1: Selected Performance Parameters (Morrissey et al. 2005) Deep Imaging Survey (DIS): The DIS consists of 20 orbit (30 ksec, mAB ' 25) exposures, over 80 deg2 , located in regions where major multiwavelength efforts are already underway. DIS regions have low extinction, low zodiacal and diffuse galactic backgrounds, contiguous pointings of 10 deg2 to obtain large cosmic volumes, and minimal bright stars. An Ultra DIS of 200 ksec, mAB ∼ 26 mag is also in progress in four fields. Nearby Galaxies Survey (NGS): The NGS targets well-resolved nearby galaxies for 1-2 eclipses. Surface brightness limits are mAB ∼27.5 mag arcsec−2 . The 200 targets are a diverse selection of galaxy types and environments (see Fig.2.3). Spectroscopic Surveys. The suite of spectroscopic surveys includes: the Wide-field Spectroscopic Survey (WSS), which covers the full 80 deg2 DIS footprint with comparable exposure time (30 ksec), and reaches mAB ∼ 20 mag for S/N∼10 spectra; the Medium Spectroscopic Survey (MSS), which covers the high priority central field in each DIS survey region (total 8 deg2 ) to mAB =21.5-23.0 mag, using 300 ksec exposures; and the Deep Spectroscopic Survey (DSS) covering 2 deg2 with 1,000 eclipses, to a depth o f mAB =23-24 mag. 2.1.2 Data collection mode GALEX performs its surveys with plans that employ a simple operational scheme requiring only two observational modes and two instrument configurations. Each orbit GALEX collects data during night segments (eclipses) and visits to a single 18 2. GALEX & GOLDMINE: A multiwavelength window on the Local Universe pre-programmed target. Each target consists either of a single pointing (single visit observation) or multiple adjacent pointings (sub-visit observations). Currently subvisits are only used for all-sky imaging survey (AIS) and in-flight calibration observations. After removing instrument overhead, each eclipse typically yields up to 1700 seconds of usable science data. During any visit or sub-visit observation the spacecraft attitude is controlled in a tight, spiraled dither. A spiral dither is used to prevent ”burn-in” of the detector active area by bright objects and to average over high spatial frequency response variations. For each sub-visit the spiral dither pattern is restarted. Since celestial sources will move on the detector, the pipeline software will reposition the time-tagged photons to common sky coordinates based on the satellite aspect solution. As many as 12 sub-visits are allowed per eclipse period (typical for AIS), with all-sky survey sub-visits obtaining 100-110 s exposure time per leg. For plans with sub-visit targets, a 20 second slew time is required to move between each leg of the observation. For some survey plans (e.g. deep imaging, spectroscopy), a single visit is insufficient to build up the requisite signal-to-noise, so a series of visits are needed in order to obtain the minimum required exposure time. 2.1.3 Counts vs. magnitudes and fluxes conversions All GALEX data are normalized to their relative exposure time, thus each count (cnt) measured on a GALEX image is in reality a cnt per sec (CPS). Below are given some equations useful to convert galaxies counts into fluxes or magnitudes. To convert from GALEX counts per sec (cps) to flux (erg cm−2 s−1 Å−1 ): F U V : F lux [erg cm−2 s−1 Å−1 ] = 1.40 × 10−15 × CPS (2.1) N U V : F lux [erg cm−2 s−1 Å−1 ] = 2.06 × 10−16 × CPS (2.2) To convert from GALEX counts per sec to magnitudes in the AB system (Oke 1974): F U V : m(AB) = −2.5 × log(CP S) + 18.82 (2.3) N U V : m(AB) = −2.5 × log(CP S) + 20.08 (2.4) Thus to convert from flux to AB magnitudes: F lux [erg cm−2 s−1 Å−1 ] F U V : m(AB) = −2.5 × log + 18.82 1.40 × 10−15 (2.5) F lux [erg cm−2 s−1 Å−1 ] + 20.08 (2.6) 2.06 × 10−16 The current estimates are that the zero-points defined here are accurate to within +/- 10% (1 sigma). N U V : m(AB) = −2.5 × log 2.1. The Galaxy Evolution Explorer 19 Figure 2.2: The transmittance profile for the NUV and FUV GALEX filters. Different galaxy spectral energy distributions are superposed. Figure 2.3: Example of GALEX image. GALEX NGS observation of NGC4631. In the color table, red-green (gold) is used for NUV, and blue for FUV. 20 2.2 2. GALEX & GOLDMINE: A multiwavelength window on the Local Universe The Galaxy On Line Database Milano Network The Galaxy On Line Database Milano Network (http://goldmine.mib.infn.it) is designed to provide access to all the data collected by G.Gavazzi, A.Boselli (Tutor and Co-Tutor of this thesis) and collaborators during several observational campaigns, started in 1985 and still in progress, aimed at providing the phenomenology of local galaxies in the widest possible frequency range. The creation of the World Wide Web site and of the MySQL database has been performed by P. Franzetti and A. Donati and a detailed description of the database architecture can be found in Donati (2004). GOLDmine is focused on 9 local clusters of galaxies: A262 (Perseus-Pisces), Cancer, A1367, A1656 (Coma), Virgo, A2147, A2151, A2197, A2199 (Hercules). In addition it contains a filament of nearly isolated galaxies, the so called “Great Wall”, thus providing the ideal laboratory for comparative analyses of galaxies in different environments, spanning a factor of 20-100 in local galaxy density. Objects are selected in the above regions with strictly optical completeness criteria. Galaxies brighter than mp = 15.7 are taken from the Catalogue of Galaxies and of Clusters of Galaxies (CGCG) by Zwicky et al. (1961) in all clusters except Virgo where objects brighter than mp = 20.0 are taken from the Virgo Cluster Catalogue (VCC) by Binggeli et al. (1985). Obviously, due to the factor of ∼ 5 difference in distance between Virgo and the other clusters, this selection limit results in dwarf galaxies being included in our database only for the Virgo cluster. However globally GOLDmine covers the whole range (4 orders of magnitude) of luminosities spanned by real galaxies. GOLDmine contains 3649 galaxies. Extensive campaigns were carried out to observe as many as possible of the 3649 target galaxies through all possible observational windows, a task that we did not complete yet. The parameters listed in the GOLDmine database are divided into 5 categories: General, Continuum and Line photometry, Dynamical and Structural. They can be obtained from GOLDmine by querying the database for an individual galaxy name or “by parameters”, “by near name or position” or “by available images”. In this case all galaxies in a given range of photographic magnitude, and morphological type can be selected. General parameters include Catalogue designations, (J2000) celestial coordinates, optical diameters, photographic magnitude, redshift, distance, morphological type. Continuum parameters include: UV, U, B, V, J, H, K magnitudes computed at the optical radius (25th mag arcsec−2 ) (see Gavazzi et al. 1996); IRAS 60 and 100 micron fluxes; radio continuum fluxes densities at 0.6 and 1.5 GHz. Line photometry includes: the atomic (HI) and molecular (H2 ) hydrogen mass; the Hα+[NII] line equivalent width and flux. Dynamical parameters include: the width of the HI line, with a quality flag; the width of the Hα line and the central velocity dispersion. 2.2. The Galaxy On Line Database Milano Network 21 Structural parameters include: the light concentration index (C31); the effective radius Re ; the effective surface brightness µe ; the total asymptotic magnitude. These quantities (see Scodeggio et al. 2002) are given separately for the H, V and B bands. The novelty of GOLDmine consists of its image section, where images can be downloaded in JPG and FITS format. Images include: Finding Charts from the Digitized Palomar Sky Survey for all galaxies. Broad band images obtained in the B, V, H and K bands. Narrow band images in the light of Hα and a red image of the underlying stellar continuum near Hα. RGB images. For some galaxies we combined several images to obtain “true” color pictures. Radial profiles of the light distribution as obtained on the available (B, V, H) images (see Gavazzi et al. 2000). When at least two radial profiles are available the color radial profile is also shown. Optical spectra integrated over the whole surface of the galaxy, obtained in drift-scan mode, i.e. by drifting the spectrograph slit over the galaxy extension (see Gavazzi et al. 2002a, 2004). Spectral Energy Distributions (SEDs) from the UV to the centimetric radio continuum obtained from broad-band photometry. The plotted data are total fluxes (extrapolated to the optical radii), unlike the individual aperture data given by NED. However they are given as observed, i.e. uncorrected for extinction from our Galaxy and for internal extinction (see Boselli et al. 2003a). It is our goal to provide a homogeneous set of keywords in all FITS header to characterize the data, including: effective integration time, filter, telescope, WCS parameters, photometric effective zero point. This homogenization is not yet complete. As also remarked in Chapter 7, the high quality of data available through GOLDMine, make this datasample one of the most appropriate for studying the evolution of nearby galaxies. Chapter 3 The FAUST-FOCA UV luminosity function of nearby clusters 3.1 Introduction The study of the galaxy luminosity function (hereafter LF) provides us with a fundamental tool for testing theories of galaxy formation and for reconstructing their evolution to the present. Recent determinations of the galaxy LF at various frequencies, in various environments (i.e.De Propris et al. 2003; Madgwick et al. 2002) and in a number of redshift intervals (i.e.Ilbert et al. 2004) have improved our knowledge of galaxy evolution and the role played by the environment in regulating the star formation activity of galaxies. The optical cluster LF is significantly steeper than that in the field (Trentham et al. 2005). This steepening is due to quiescent galaxies, more frequent at low luminosities in clusters, while the LF of cluster star forming objects is similar to that in the field (De Propris et al. 2003). The causes of this difference might reside in the density-morphology relation (Dressler 1980; Whitmore et al. 1993) and in particular in the overabundance of dwarf ellipticals in rich clusters (Ferguson & Sandage 1991), whose origin is currently debated in the framework of the environmental effects on galaxy evolution. The ultraviolet emission UV( ∼ 2000 Å), being dominated by young stars of intermediate masses (2 < M < 5 M , Boselli et al. 2001) represents an appropriate tool to identify and quantify star formation activity. Although before the launch of GALEX, the shape of local field UV LF (Sullivan et al. 2000) was supposed to be well determined, there was still a fair amount of uncertainty on the UV luminosity function of clusters. Its slope was undetermined due to the insufficient knowledge of the background counts (Cortese et al. 2003b). Andreon (1999) proposed a very steep faint end (α ∼ −2.0, −2.2), significantly different from the field LF (α ∼ −1.5). However Cortese et al. (2003b) pointed out that this steep slope is likely caused by an un23 24 3. The FAUST-FOCA UV luminosity function of nearby clusters derestimation of the density of background galaxies and proposed a flatter faint-end slope (α ∼ −1.35 ± 0.20). Unfortunately the statistical uncertainty was too high for making reliable comparisons between the cluster and the field LFs. In this chapter I re-compute the cluster UV luminosity function with two major improvements over previous determinations. We increase the redshift completeness of the UV selected sample using new spectroscopic observations of Coma and Abell 1367 (see Chapter 5 and Cortese et al. 2004), and compute for the first time the UV LF of the Virgo cluster. These improvements are not sufficient to constrain the LF of each individual cluster, however the UV composite luminosity function, constructed for the first time in this paper can be compared with that of the field. Doing so I try anticipating one of the main goals of the Galaxy Evolution Explorer (GALEX) which, as shown in the next Chapter, will help us shade light on the UV properties of galaxies and their environmental dependences. We assume a distance modulus µ= 31.15 for the Virgo cluster (Gavazzi et al. 1999a), µ=34.80 for Abell 1367 and µ=34.91 for the Coma cluster (Gavazzi et al. 1999b). 3.2 The Data The sample analyzed in this chapter comprises the UV sources detected in Virgo, Coma and Abell 1367 clusters by the FOCA (Milliard et al. 1991) and FAUST (Lampton et al. 1990) experiments. The FOCA balloon-borne wide field UV camera (λ = 2000Å; ∆λ = 150Å) observed ∼ 3 square degrees (∼ 8 Mpc2 ) in the Abell 1367 (unpublished data) and Coma clusters (Donas et al. 1995) and ∼ 12 square degrees (∼ 1 Mpc2 ) in the Virgo cluster (data are taken from the extragalactic database GOLDMine, Gavazzi et al. 2003a). The FOCA observations of Virgo are not sufficient to compile a complete catalog: no sources brighter than mUV ∼ 12.2 were detected due to the small area covered. We thus complement the UV database with the wide field observations performed by the FAUST space experiment (λ = 1650Å; ∆λ = 250Å) in the Virgo direction (Deharveng et al. 1994), covering ∼ 100 square degrees (∼ 8.8 Mpc2 ). The FAUST completeness limit is mUV ∼ 12.2 (Cohen et al. 1994), significantly lower than the FOCA magnitude limit: mUV ∼ 18.5. However combining the two UV catalogs we hope to constrain the shape of the UV luminosity function across 7 magnitudes. We use the FAUST observations for mUV < 12.2 and the FOCA observations for mUV ≥ 12.2. To account for the different response function of FAUST and FOCA filters we transform the UV magnitudes taken by FAUST at 1650Å assuming a constant color index: UV(2000) = UV(1650) + 0.2 mag (Deharveng et al. 1994, 2002). We think however that this difference does not bias the galaxy populations selected by the two experiments. The estimated error on the UV magnitudes is 0.3 mag in general, but it ranges from 0.2 mag for bright galaxies, to 0.5 mag for faint sources observed in frames with larger than average calibration uncertain- 3.3. The UV luminosity functions 25 ties. The UV emission associated with bright galaxies is generally clumpy, thus it has been obtained by integrating the flux over the galaxy optical extension, determined at the surface brightness of 25 mag arcsec−2 in the B-band. The spatial resolution of the UV observations is 20 arcsec and 4 arcmin for FOCA and FAUST respectively. The astrometric accuracy is therefore insufficient for unambiguously discriminating between stars and galaxies. To overcome this limitation, we cross-correlate the UV catalogs with the deepest optical catalogs of galaxies available: the Virgo Cluster Catalog (VCC, Binggeli et al. 1985), complete to mB ∼ 18, for the Virgo cluster and the r 0 band catalog by Iglesias-Páramo et al. (2003), complete to mr0 ∼ 20, for Coma and Abell 1367. We used as matching radius the spatial resolution of each observation. In case of multiple identifications we select the galaxy closest to the UV position. The resultant UV selected sample is composed of 156 galaxies in Virgo, 140 galaxies in Coma and 133 galaxies in Abell 1367. 3.3 The UV luminosity functions Unlike the VCC catalog, the Coma and A1367 r 0 catalogs used for star/galaxy discrimination do not cover all the area observed by FOCA but only the cluster cores. This reduces our analysis to an area of ∼ 1 square degrees (∼ 2.6 Mpc2 ) in Coma and ∼ 0.7 square degrees (∼ 1.8 Mpc2 ) in Abell 1367. Including new spectroscopic observations (Cortese et al. 2004), the redshift completeness of the UV selected sample reaches the 65% in Abell 1367, the 79% in Coma and the 83% in Virgo. The redshift completeness per bin of magnitude of each cluster is listed in Table 3.1. We remark that for MUV ≤ −16.5 (corresponding to the FOCA magnitude limit in Coma and Abell1367), the redshift completeness of the Virgo cluster sample is 98%. As discussed by Cortese et al. (2003b), the general UV galaxy counts (Milliard et al. 1992) are uncertain and cannot be used to obtain a reliable subtraction of the background contribution from the cluster counts. Therefore, in order to compute the cluster LF, we use the statistical approach recently proposed by De Propris et al. (2003) and Mobasher et al. (2003). We assume that the UV spectroscopic sample is ’representative’, in the sense that the fraction of galaxies that are cluster members is the same in the (incomplete) spectroscopic sample as in the (complete) photometric sample. For each magnitude bin i we count the number of cluster members NM , the number of galaxies with a measured recessional velocity NZ and the total number of galaxies NT . The completeness-corrected number of cluster members in each bin is: Ni = NM NT NZ (3.1) 26 3. The FAUST-FOCA UV luminosity function of nearby clusters Table 3.1: Integral redshift completeness in bin of 0.5 magnitudes. MUV ≤ −21.75 −21.25 −20.75 −20.25 −19.75 −19.25 −18.75 −18.25 −17.75 −17.25 −16.75 Redshift completeness Virgo Coma Abell1367 − − 100% 100% 92% 95% 97% 97% 97% 98% 98% − 100% 100% 100% 100% 100% 100% 97% 95% 84% 79% 100% 100% 100% 100% 100% 100% 100% 100% 95% 80% 65% NT is a Poisson variable, and NM is a binomial variable (the number of successes in NZ trials with probability NM /NZ ). Therefore the errors associated with Ni are given by: δ 2 Ni 1 1 1 = + − (3.2) 2 Ni NT NM NZ The completeness-corrected number of cluster members obtained from (3.1) are given in Table 3.2 and the luminosity functions for the four studied samples are shown in Fig.3.1. The two different datasets used for the Virgo cluster have only one magnitude bin (MUV = −18.75) overlap. In this bin the two LFs are in agreement and there is no indication that a change of slope occurs. We thus feel comfortable combining them into a composite Virgo UV luminosity function across 7 magnitudes. In order to determine whether the LFs of the three clusters are in agreement we perform a two-sample χ2 test. We obtain P (χ2 ≥ χ2obs ) ∼82% for the Virgo and Abell1367 LFs, P (χ2 ≥ χ2obs ) ∼87% for the Virgo and the Coma cluster LFs and P (χ2 ≥ χ2obs ) ∼98% for the Coma and Abell1367 LFs, pointing out that the three LFs are in fair agreement within their completeness limits. 3.3.1 The composite cluster luminosity function The uncertainties of each individual cluster luminosity function are too large to fit a complete Schechter (Schechter 1976) function to the data and compare it with the field UV LF. However combining the three data-sets analyzed in this paper we 3.3. The UV luminosity functions 27 Figure 3.1: The UV luminosity functions for the four analyzed data sets. compute the UV composite luminosity function of 3 nearby clusters. The composite LF is obtained following Colless (1989), by summing galaxies in absolute magnitude bins and scaling by the area covered in each cluster. The number of galaxies in the jth absolute magnitude bin of the composite LF (Ncj ) is given by: Ncj = 1 X Nij mj i A i (3.3) where Nij is the completeness-corrected number of galaxies in the jth bin of the ith cluster, Ai is the area surveyed in the ith cluster and mj is the number of clusters contributing to the jth bin. The errors in Nij are computed according to: δNcj = 1 h X δNij 2 i1/2 mj i Ai (3.4) where δNij is the error in the jth bin of the ith cluster determined in (3.2). The weight associated to each cluster is computed according to the surveyed area, instead of the number of galaxies brighter than a given magnitude, as used by Colless (1989). The UV composite luminosity function is given in Fig.3.2 in the full magnitude range. However since for magnitudes fainter than MUV ∼ −16.5 the only available data are the Virgo FOCA observations, we fit the composite luminosity function with 28 3. The FAUST-FOCA UV luminosity function of nearby clusters Figure 3.2: The composite UV luminosity function of 3 nearby clusters. The solid line represents the best Schechter fit to the data for MUV ≤ −16.5. the Schechter functional form (Schechter 1976): φ(MUV ) = 0.4 ln 10 φ∗ 100.4(M ∗ −M UV )(α+1) e−10 0.4(M ∗ −MUV ) only for MUV ≤ −16.5, that is the completeness limit in Coma and Abell 1367. ∗ The resulting Schechter parameters are MUV = −20.75 ± 0.40 and α = −1.50 ± 0.10. The faint-end slope is consistent within 1 σ with the lower limit for Coma and A1367 recently proposed by Cortese et al. (2003b), but significantly flatter than the slope α ∼ −2.0, −2.2 found for Coma by Andreon (1999), suggesting that this very steep luminosity function was due to an underestimate of the density of background galaxies. 3.4 Discussion Although the UV(2000 Å) radiation is dominated by young stars of intermediate masses (2<M<5M , Boselli et al. 2001), it is frequently detected also in early-type galaxies with no recent star formation episodes (Deharveng et al. 2002). Unfortunately we have no morphological (or spectral) classification for all the UV selected galaxies in order to separate the contribution of late and early type galaxies. How- 3.4. Discussion 29 Figure 3.3: The UV bi-variate composite luminosity functions of nearby clusters. Red (UV − B > 2) and blue (UV − B < 2) galaxies are indicated with empty and filled circles respectively. ever, based on the spectral energy distributions computed by Gavazzi et al. (2002a), we can use the total color UV − B, available for the 94% of galaxies in our sample, to discriminate between red elliptical (UV − B > 2) and blue spiral (UV − B < 2) galaxies. B magnitudes are taken from the VCC (Binggeli et al. 1985), the Godwin et al. (1983) catalog and the Godwin & Peach (1982) catalog for Virgo, Coma and Abell 1367 respectively. The bi-variate composite luminosity function derived for galaxies of known UV − B color is shown in Fig.3.3. It shows that the star forming galaxies dominate the UV LF for MUV ≤ −18, as Donas et al. (1991) concluded for the first time. Conversely, for MUV ≥ −17.5, the number of red and blue galaxies is approximately the same, pointing out that, at low luminosities, the UV emission must be ascribed not only to star formation episodes but also to Post-Asymptotic Giant Branch (PAGB) low mass stars in early type galaxies (Deharveng et al. 2002). Similarly, if we restrict the analysis to the fraction (∼ 50 %) of objects with known morphological type, we find that late-types (Sa or later) dominate at bright UV luminosities, while early-type objects contribute at the faint UV levels. Since Virgo and Abell1367 are spiral-rich clusters while Coma is spiral-poor, one might expect that the LFs of the three clusters obtained combining all types should have different shapes, contrary to the observations. The point is that the combined LF of the two types is dominated, at high UV lumi- 30 3. The FAUST-FOCA UV luminosity function of nearby clusters Table 3.2: The completeness-corrected differential number of galaxies per bin of magnitude MUV mag −21.75 −21.25 −20.75 −20.25 −19.75 −19.25 −18.75 −18.25 −17.75 −17.25 −16.75 Ni Virgo Virgo Coma (Faust) (Foca) 0 0 2 1 7 9 13 0 0 0 0 0 0 0 0 0 0 2 2 3 3 4 0 1 0 5 3 3 5 8.6 7.7 15.8 18.6 Abell 1367 1 0 1 1 4 4 3 6 6.7 10.1 12.7 nosity by the spiral component, while at low luminosity early- and late-type galaxies contribute similarly. The UV LF of the spiral component are similar in the three clusters. At faint UV luminosities also the number density of early-type galaxies is approximately the same in the three clusters. Only at relatively high UV luminosity the number density of early-type galaxies in the Coma cluster exceeds significantly that of the other two clusters, but it is still much lower than the one of the late-type component. Therefore the LF obtained by combining early- with late-type galaxies results approximately the same in the three clusters. The cluster composite luminosity function has identical slope and similar M ∗ as the ∗ UV luminosity function computed by Sullivan et al. (2000) for the field: MUV = −21.21 ± 0.13, α = −1.51 ± 0.10, as shown in Fig. 3.4. This result is quite surprising since we have just shown that at low luminosity the contribution of ellipticals is not negligible, and early-type galaxies are expected to be more frequent in high density environments. This result seems in contradiction with recent studies of cluster galaxies carried out in Hα (Iglesias-Páramo et al. 2002) and B-bands (De Propris et al. 2003). They find that the LFs of star forming galaxies in clusters and in the field have the same shape, contrary to early type galaxies in clusters that have a brighter and steeper LF than their field counterparts (De Propris et al. 2003). In order to understand this apparent difference between optical and UV luminosity functions we needed to wait the launch of GALEX and higher quality (and more homogeneous) 3.4. Discussion 31 Figure 3.4: The cluster and the field UV luminosity functions. The composite cluster LF is given with filled circles. The solid line indicates the best Schechter fit of the field LF of Sullivan et al. (2000). The normalization is such that the two LFs match at MUV ∼ −19.25. UV observations. Chapter 4 GALEX UV luminosity function of Abell1367 4.1 Introduction As I have shown in the previous Chapter, before the launch of the Galaxy Evolution Explorer (GALEX), the FOCA experiment allowed the first determinations of the UV LF of local field galaxies (Treyer et al. 1998; Sullivan et al. 2000) and of nearby clusters (Donas et al. 1991; Andreon 1999; Cortese et al. 2003b). Combining the FOCA and FAUST data Cortese et al. (2003a) determined the first composite LF of nearby clusters. They found no significant differences with the LF in the field. However this early determination was affected by large statistical errors due to the uncertainty in the UV background counts (Cortese et al. 2003b). GALEX has opened a new era of extragalactic UV astronomy. In particular it provides for the first time precise UV photometry of galaxies over large stretches of the sky (Xu et al. 2005), thus making the background subtraction method more reliable than in the past. Moreover its higher sensitiveness, higher resolution, large field of view make GALEX observations a unique homogeneous sample for statistical analysis of galaxies UV properties. 4.2 UV data GALEX provides far-ultraviolet (FUV; λeff = 1528Å, ∆λ = 442Å) and near-ultraviolet (NUV; λeff = 2271Å, ∆λ = 1060Å) images with a circular field of view of ∼ 0.6 degrees radius. The spatial resolution is ∼5 arcsec. The data analyzed in this Chapter consist of two GALEX pointings of the Abell cluster 1367, with a mean exposure time of 1460s, , centered at R.A.(J2000)=11:43:41.34 Dec(J.2000)=+20:11:24.0 (e.g. offset to the north of the cluster to avoid a star bright enough to threaten the detector, see Fig.4.1). Sources were detected and measured using SExtractor (Bertin & Arnouts 33 34 4. GALEX UV luminosity function of Abell1367 Figure 4.1: The GALEX observation of Abell1367. ROSAT X-ray contour are superposed in black. The tick rectangular region indicates the region covered by the optical catalogues used for the star/galaxy discrimination. 1996). The 100% completeness limit is mAB ∼ 21.5 both in FUV and NUV (Xu et al. 2005). As the NUV images are significantly deeper than the FUV, sources were selected and their parameters determined in the NUV. FUV parameters were extracted in the same apertures. We used a larger SExtractor deblending parameter compared to the standard GALEX pipeline, providing reliable MAGAUTO also for very extended sources. The calibration uncertainty of the NUV and FUV magnitudes is ∼ 10% (Morrissey et al. 2005). Magnitudes are corrected for Galactic extinction using the Schlegel et al. (1998) reddening map and the Galactic extinction curve of Cardelli et al. (1989). The applied extinction corrections are of 0.18 and 0.17 mag for the NUV and FUV bands respectively. To avoid artifacts present at the edge of the 4.2. UV data 35 Figure 4.2: Comparison between FOCA (upper image) and GALEX (lower image) observation of the center of Abell1367. It emerges clearly the strong improvement in resolution and sensitiveness of new GALEX data. 36 4. GALEX UV luminosity function of Abell1367 Figure 4.3: Left: The comparison between FOCA and GALEX NUV (left) and FUV (right) magnitudes of galaxies in Abell1367. The continuum line shows the best linear fit to the data. field, we considered only the central 0.58 deg radius from the field center. A reliable star/galaxy discrimination was achieved by matching the GALEX catalog against the deepest optical catalogs available for A1367 (B < 22.5 mag and r 0 < 21 mag; Iglesias-Páramo et al. 2003), using a search radius of 6 arcsec, as adopted by Wyder et al. (2005) for the estimate of the GALEX local field LF. The optical catalogs do not include a negligible part (∼ 0.09 square degrees) of the GALEX field. A total number of 292 galaxies in the FUV and of 480 galaxies in NUV with mAB ≤ 21.5 are detected in the ∼ 0.96 square degrees field (∼ 2.5Mpc2 ) analyzed in this Chapter. Great part of the field observed by GALEX covers the area studied in the previous Chapter with FOCA observations. The two observations of the cluster center are presented in Fig.4.2: emerges clearly the strong improvement in resolution and sensitiveness of new GALEX data. In Fig.4.3 (left) we compared the UV magnitudes determined from FOCA and from GALEX NUV observations for the 96 galaxies detected by both instruments. The two sets of measurements are in satisfactory agreement. The linear regression between the two datasets is: MGALEX (2310Å) = (1.02 ± 0.03) × MFOCA (2000Å) + (1.74 ± 0.51) (4.1) MGALEX (1530Å) = (1.04 ± 0.04) × MFOCA (2000Å) + (1.71 ± 0.70) (4.2) with a mean dispersion of 0.23 and 0.32 mag in NUV and FUV bands respectively, consistent with the error assumed in the previous Chapter for FOCA observations. 4.3. The luminosity function 37 Figure 4.4: The redshift completeness per bin of UV magnitude in Abell 1367. Band NUV NUV FUV FUV UV(2000Å) Sample A1367 F ield A1367 F ield Composite cluster Schechter P arameters M∗ α −19.77 ± 0.42 −18.23 ± 0.11 −19.86 ± 0.50 −18.04 ± 0.11 −18.79 ± 0.40 −1.64 ± 0.21 −1.16 ± 0.07 −1.56 ± 0.19 −1.22 ± 0.07 −1.50 ± 0.10 Table 4.1: Best Fitting Parameters. 4.3 The luminosity function The determination of the cluster LF requires a reliable estimate of the contribution from background/foreground objects to the UV counts. This can be accurately achieved for mAB ≤ 18.5, since at this limit our redshift completeness is ∼ 90 % (Cortese et al. 2003b, 2004; see Fig. 4.4). The redshift completeness drops rapidly at magnitudes fainter than mAB ∼ 18.5, thus requiring the contamination to be estimated statistically. Two methods are usually applied for the computation of cluster LFs. The first one is based on the statistical subtraction of field galaxies, per bin of UV magnitude, that are expected to be randomly projected onto the cluster area, as derived by Xu et al. (2005). Alternatively, the completeness corrected method proposed by De Propris et al. (2003) is to be preferred when the field counts have large uncertainties. It is based on the assumption that the UV spectroscopic sample 38 4. GALEX UV luminosity function of Abell1367 Figure 4.5: The GALEX NUV (left) and FUV (right) LF for Abell 1367. Open dots are obtained using the subtraction of field counts obtained by Xu et al. (2005); filled dots are obtained using the completeness corrected method. The solid line represents the best Schechter fit. The dotted line shows the composite nearby clusters 2000 Å LF by (Cortese et al. 2003a). The dashed line represents the GALEX local field LF (Wyder et al. 2005), normalized in order to match the cluster LF at MAB ∼ −17.80. Figure 4.6: The NUV (left) and FUV (right) bi-variate LF of A1367. Star-forming and quiescent galaxies are indicated with empty triangles and filled squares respectively. The dashed line represents the GALEX local field LF (Wyder et al. 2005), normalized as in Fig.4.5 4.3. The luminosity function 39 (e.g. membership confirmed spectroscopically) is ’representative’ of the entire cluster, i.e. the fraction of galaxies that are cluster members is the same in the (incomplete) spectroscopic sample as in the (complete) photometric one. For each magnitude bin i we count the number of cluster members NM (i.e. galaxies with velocity in the range 4000<V<10000 km s−1 ; Cortese et al. 2004), the number of galaxies with a measured recessional velocity NZ and the total number of galaxies NT . The ratio NZ /NT , corresponding to the redshift completeness in each magnitude bin is shown in Fig.4.4. The completeness-corrected number of cluster members in each bin is Ni = (NM × NT )/NZ . NT is a Poisson variable, and NM is a binomial variable (the number of successes in NZ trials with probability NM /NZ ). Therefore the errors associated with Ni are given by (δ 2 Ni /Ni2 ) = (1/NT ) + (1/NM ) − (1/NZ ). The NUV and FUV LFs using both methods (see Fig 4.5) are in good agreement for MAB ≥ −14.3. In the last bin the two methods are inconsistent as the completeness corrected method predicts a higher slope than the statistical background subtraction. This disagreement is likely due to the severe redshift incompleteness for MAB ≥ −14.3. In any case we take the weighted mean of the two determinations. Due to the small number of galaxies populating the high luminosity bins (i.e. only three objects brighter than MAB ∼ −18.3), the LFs are not well fitted with a Schechter function (Schechter 1976): the best-fit M∗ turns out to be brighter than the brightest observed galaxy.For this reason we first determine the faint-end (−18.3 ≤ MAB ≤ −13.3) slope in each band, fitting the LFs with a power law (Φ(M ) = c 10kM ) by minimizing χ2 . The α parameter of the Schechter function can be derived from k using the relation α = −(k/0.4 + 1). Then we fit the LFs with a Schechter function, keeping α fixed to the value previously obtained. This is not the canonical Schechter fit, but it provides a more realistic set of parameters than using a three-free-parameter fit. The best fit parameters and their errors are listed in Table.4.3. In order to separate the contribution to the LF of star-forming from quiescent galaxies, we divide the sample into two classes. Using Hα imaging data (Iglesias-Páramo et al. 2002; Gavazzi et al. 1998, 2002b, 2003a) and optical spectroscopy (Cortese et al. 2003b, 2004) we can discriminate between star-forming (EW (Hα) > 0 Å) and quiescent (EW (Hα) = 0 Å) objects. Unfortunately neither UV field counts for different morphological types nor a measure of EW (Hα) for all the UV selected galaxies are available. Thus we can only apply the completeness corrected method to determine the bi-variate LFs. We assume that in each bin of magnitude the fraction of starforming and quiescent cluster members is the same in the (incomplete) spectroscopic sample as in the (complete) photometric sample. The bi-variate LFs derived by this method are shown in Fig.4.6. 40 4. GALEX UV luminosity function of Abell1367 4.4 Discussion As shown in Fig.4.5, the GALEX LFs have a shape consistent with the composite LF of nearby clusters as constructed in the previous Chapter (see also Cortese et al. (2003a)). Conversely, whatever fitting procedure one adopts, they show a steeper faint-end slope and a brighter M ∗ than the GALEX field LF recently determined by Wyder et al. (2005). In fact the GALEX local field luminosity function shows a faintest bright-end and a flatter faint end than the previous determination by Sullivan et al. (2000), but the reason for this difference is not yet clearly understood. Wyder et al. (2005) argued that magnitudes estimated by FOCA are on average brighter than the GALEX one, with the difference becoming larger for fainter sources; suggesting that these offsets and nonlinearities in the FOCA photometry could account for a major part of the observed difference between the two field luminosity functions. However we have shown that this seems not the case at least for Abell1367 observations. On the contrary I think that part of the problem could be due not to different photometric estimates but to the different areas used by GALEX and FOCA to estimate the field LF. In the case of FOCA, Treyer et al. (1998) and Sullivan et al. (2000) used the pointing of Abell 1367 to estimate the field LF: thus a partial contamination of cluster galaxies could explain why the FOCA field and cluster LF results very similar. The brighter M ∗ observed in Abell1367 is probably to be ascribed to the particular galaxy population of this cluster. In fact Abell 1367 is a young cluster of galaxies composed of at least four dynamical units at the early stage of a multiple merging event (see Chapter 5 and Cortese et al. 2004). Some galaxies have their star formation enhanced due to interaction with the cluster environment, and it is this population that is responsible for the bright M ∗ observed in this cluster. Conversely the high faint-end slope observed in this cluster is due to the significant contribution of non star-forming systems at faint UV magnitudes. In fact, as shown in Fig.4.6, star-forming galaxies dominate the UV LF for MAB ≤ −17 mag, as Donas et al. (1991) concluded for the first time. For MAB ≥ −16 mag however, the number of red galaxies increases very rapidly1 . This result is consistent with an UV LF constructed starting from the r 0 LF computed by Iglesias-Páramo et al. (2003): assuming a mean color N U V − r 0 ∼ 1 mag and N U V − r 0 ∼ 5 for star-forming and quiescent galaxies respectively, we are able to reproduce the contribution, at low UV luminosities, of elliptical galaxies. Moreover the difference observed between NUV and FUV cluster LFs can be understood looking at the FUV-NUV color magnitude relation (computed only for confirmed cluster members) shown in Fig.4.7. The star-forming objects dominate at high UV luminosities while the quiescent systems contribute more 1 The bi-variate LFs cannot be compared with the ones computed by Treyer et al. (2005) for the field, since their samples do not contain ellipticals but only spiral galaxies. 4.4. Discussion 41 Figure 4.7: The FUV-NUV color magnitude relation for confirmed members of A1367. Symbols are as in Fig.4.6 Figure 4.8: The optical (r 0 -band) distribution for star forming (blue histogram) and quiescent (red histogram) galaxies in our sample. 42 4. GALEX UV luminosity function of Abell1367 at faint magnitudes. Their mean FUV-NUV color is ∼ 1.5 mag thus they influence the LF at higher luminosities in the NUV than in the FUV (see Fig.4.6). The optical luminosity distribution of star forming and quiescent systems, presented in Fig. 4.8, points out clearly that early type galaxies contributing to the UV faint end slope are the giant, optically bright, galaxies that dominate the bright end of the optical luminosity functions. This means that, in UV, the steeper faint end slope observed in clusters is only due to the contribution of giant ellipticals and not of dwarf elliptical galaxies, as observed at optical wavelengths. We can thus conclude that, in clusters, a significant fraction of the low luminosity UV emission comes massive early type galaxies. This result is expected since in the field the fraction of quiescent systems is significantly lower than that of star forming objects (Dressler 1980; Whitmore et al. 1993), thus their contribution to the LF is negligible. Moreover, the UV emission of ellipticals has a different nature from the one emitted by star forming systems. In fact it does not arise from newly born stars but from low mass old post asymptotic giant branch stars (O’Connell 1999), as I will discuss in depth in Chapter 6. Finally, the LFs of cluster star-forming systems have a faint-end slope (α ∼ −1.25 ± 0.2) consistent within the statistical uncertainties with the GALEX field LF. The similar shape observed in the LF of star forming galaxies in different environments goes in the same direction with recent studies of cluster galaxies carried out in Hα (Iglesias-Páramo et al. 2002) and B-bands (De Propris et al. 2003). They find that the LFs of star forming galaxies in clusters and in the field have the same shape, contrary to early type galaxies in clusters that have a brighter and steeper LF than their field counterparts (De Propris et al. 2003). This indicates that, whatever mechanism (i.e. ram pressure, tidal interaction, galaxy harassment) quenches/enhances the star formation activity in late-type cluster galaxies, it influences similarly and with a short time scale the giant and the dwarf components , so that the shape of their LF is unchanged and only the normalization is modified. Chapter 5 Multiple merging in Abell1367 5.1 Introduction Clusters of galaxies represent the most massive gravitationally bound systems in the Universe. They provide us with valuable insights into the formation of large-scale structures, as well as into the formation and evolution of galaxies. The hierarchical model predicts that galaxy clusters are formed by accretion of units of smaller mass at the nodes of large-scale filaments (West et al. 1991; Katz & White 1993). Statistical analyses of clusters have shown that even at low redshift a high fraction of clusters presents substructures, implying that clusters are still dynamically young units, undergoing the process of formation (Dressler & Shectman 1988). The Abell cluster 1367 (z ∼ 0.0216) lies at the intersection of two filaments, the first extending roughly 100 Mpc from Abell 1367 toward Virgo (West & Blakeslee 2000), the second connecting Abell 1367 to Coma (as a part of the Great Wall, Zabludoff et al. 1993). With its irregular X-ray distribution (Jones et al. 1979; Bechtold et al. 1983; Grebenev et al. 1995), high fraction of spiral galaxies and low central galaxy density, Abell 1367 can be considered as the prototype of a nearby dynamically young cluster. ASCA X-ray observations pointed out the existence of a strong localized shock in the intra-cluster medium (ICM) suggesting that Abell 1367 is experiencing a merging between two substructures (Donnelly et al. 1998). Moreover recent Chandra observations (Sun & Murray 2002), and a preliminary analysis of the XMM data (Forman et al. 2003), indicate the presence of cool gas streaming into the cluster core, supporting a multiple merger scenario. Optical and radio observations also suggest that this cluster is currently experiencing galaxy infall into its center. Gavazzi et al. (1995, 2001a) discovered two head-tail radio sources associated with disk galaxies with an excess of giant HII regions on their leading edges, in the direction of the cluster center. The observational scenario 43 44 5. Multiple merging in Abell1367 is consistent with the idea that ram-pressure (Gunn & Gott 1972) is, for a limited amount of time, enhancing the star formation of galaxies that are entering the cluster medium. In addition Gavazzi et al. (2003b) pointed out the existence of a group of star bursting galaxies infalling into the cluster core. Although X-ray, radio and optical observations suggest that Abell 1367 is dynamically young and it is still undergoing the process of formation, detailed spatial and dynamical analysis of this cluster has not been attempted so far. Girardi et al. (1998) detected a secondary peak in the cluster velocity distribution, suggesting that Abell 1367 is a binary cluster, but their analysis was based on ∼ 90 redshifts, insufficient for drawing a detailed model of the cluster kinematics. Cortese et al. (2003b) carried out a deep (r 0 < 20.5) spectroscopic survey of the central ∼ 1.3 square degrees of Abell 1367 adding 60 new spectra (33 members). Here I present new measurements for 119 galaxies (adding another 33 cluster members). In total 273 redshifts were measured in the region, out of which 146 are cluster members, allowing the first detailed dynamical analysis of Abell 1367. 5.2 Observations and data reduction The cluster region analyzed in this Chapter covers an area of ∼ 1.3 square degrees centered at α(J.2000) = 11h44m00s δ(J.2000) = 19d43m30s. r 0 imaging material was used to extract a catalogue of galaxy candidates in Abell 1367 complete to r 0 ∼ 20.5 mag, and to select the targets of the present spectroscopic survey. Spectroscopy of Abell 1367 was obtained with the AF2-WYFFOS multi fiber spectrograph at the 4.2m William Herschel Telescope (WHT) on La Palma (Spain) during 2003, March 27-29. WYFFOS has 150 science fibers of 1.6 arcsec diameter coupled to a benchmounted spectrograph which relies on a TEK 1024 × 1024 CCD. The 316R grating was used, giving a dispersion of ∼240 Å/mm, a resolution of ∼ 6Å FWHM, and a total spectral coverage of ∼5600 Å. The spectra were centered at ∼ 6500Å, thus covering from 3600 Å to 9400 Å. We allocated typically ∼ 70 objects to fibers in a given configuration and, on average, 15 sky fibers. A total of 4 configurations were executed, with an exposure time of 4x1800 sec for each configuration. Argon lamps for wavelength calibration were obtained for each exposure. The reduction of the multi fiber spectra was performed in the IRAF1 environment, using the IMRED package. After bias subtraction, the apertures were defined on dome flat-field frames and used to trace the spectra on the CCD. The arc spectra were extracted and matched with arc lines to determine the dispersion solution. The 1 IRAF is distributed by the National Optical Astronomy Observatories, which is operated by the Association of Universities for Research in Astronomy, Inc., under the cooperative agreement with the National Science Foundation 5.2. Observations and data reduction Observatory WHT Cananea Loiano 45 Dates N. gal. Spectrograph Dispersion Å/mm Coverage Å CCD pix µm March 03 March 03 March 03 - Feb. 04 98 12 9 AF2-WYFFOS LFOSC BFOSC 240 228 198 3600-9400 4000-7100 3600-8900 1024 × 1024 T EK 576 × 384 T H 1340 × 1300 EEV 24 23 20 Table 5.1: The spectrograph characteristics rms uncertainty of the wavelength calibration ranged between 0.1 and 0.3 Å. The lamps’ wavelength calibration was checked against known sky lines. These were found within ∼ 0.5 Å of their nominal position, providing an estimate of the systematic uncertainty on the derived velocity of ∼ 25 km s−1 . The object spectra were extracted, wavelength calibrated and normalized to their intensity in the interval 5400-5600 Å. A master sky spectrum, that was constructed by combining various sky spectra was normalized to each individual science spectrum and then subtracted from it. Unfortunately strong sky residuals were left after this procedure, limiting the number of useful spectra to 98 (as listed in Tab. 5.9). Nine additional long-slit, low dispersion spectra were obtained in March 2003 and in February 2004 using the imaging spectrograph BFOSC attached to the Cassini 1.5m telescope at Loiano (Italy). Another twelve spectra were taken with LFOSC at the 2.1m telescope of the Guillermo Haro Observatory at Cananea (Mexico). These observations were performed using a 2.0 arcsec slit and the wavelength calibration was secured with exposures of HeAr and XeNe lamps at Loiano and Cananea respectively. The on-target exposure time ranged between 15 and 30 min according to the brightness of the targets. After bias subtraction, when 3 or more frames of the same target were obtained, these were combined (after spatial alignment) using a median filter to help cosmic rays removal. Otherwise the cosmic rays were removed using the task COSMICRAYS and/or under visual inspection. The lamps wavelength calibration was checked against known sky lines. These were found within ∼ 1 Å from their nominal position, providing an estimate of the systematic uncertainty on the derived velocity of ∼ 50 km s−1 . After subtraction of sky background, one-dimensional spectra were extracted from the frames. The redshift were obtained using the IRAF FXCOR Fourier cross-correlation (Tonry & Davis 1979) task, excluding the regions of the spectra affected by night-sky lines. Moreover all the spectra and their best correlation function were visually examined to check the redshift determination. Table 5.2 lists the characteristics of the instrumentation in the adopted set-up. The 119 new velocity measurements presented in this Chapter are listed in Table 5.9 as follows: Column 1: Galaxy designation. 46 5. Multiple merging in Abell1367 Figure 5.1: Cumulative redshift distribution for galaxies in the studied region. Column 2, 3: (J2000) celestial coordinates, measured with few arcsec uncertainty. Column 4: r 0 band magnitude. Column 5: observed recessional velocity. Column 6: telescope (WHT=William Herschel Telescope; LOI=Loiano; CAN=Cananea) Combining the new set of 119 redshifts (given in Tab. 5.9) with the ones available from the literature (NED; Cortese et al. 2003b; Rines et al. 2003), we have the redshift for 273 galaxies of which 146 are cluster members (4000 km s−1 ≤ V ≤ 10000 km s−1 ). The cumulative redshift distribution, in the observed area, as a function of the r 0 magnitude is shown in Fig.5.1. The completeness is ∼ 70% at r 0 < 17.5, and it drops to ∼ 45% at r 0 < 18.5. 5.3 The global velocity distribution The line of sight (LOS) velocity distribution for the 146 cluster members is shown in Fig. 5.2. The mean and standard deviation are known to be efficient estimators of the central location and scale when the underlying population is gaussian. Unfortunately they are not minimum variance estimators when the nature of the observed population is significantly non-Gaussian. The best location and scale estimators must be resistant to the presence of outliers and robust to a broad range of non-Gaussian underlying populations. Thus, following Beers et al. (1990), we consider the biweight estimator as the best estimator of location (CBI ) and scale (SBI ) of the cluster ve- 5.3. The global velocity distribution 47 Figure 5.2: Velocity histogram and stripe density plot for the members of Abell 1367. Arrows mark the location of the most significant weighted gaps in the velocity distribution. locity distribution. We find a location CBI = 6484 ± 81 km s−1 and a scale SBI = 891 ± 58 km s−1 , in agreement with previous studies (e.g. Girardi et al. 1998; Struble & Rood 1999). Visual inspection of Fig. 5.2 suggests that the velocity distribution differs from a Gaussian, a deviation that should be quantified using appropriate statistical tests. We analyze the higher moments of the distributions using the kurtosis and the skewness shape estimators. Kurtosis indicates a difference in the tails length compared to a Gaussian (positive kurtosis is indicative of long tails). Skewness indicates asymmetry (positive skewness implies that the distribution is depleted from values lower than the mean location, conversely negative skewness denotes a depletion of values higher than the mean). In addition we calculate the asymmetry index (AI) and tail index (TI) introduced by Bird & Beers (1993) as alternatives to the distribution higher moments. These indicators measure the shape of a distribution but, contrary to skewness and kurtosis, which depend on the estimate of the location and the scale of the underlying distribution, they are based on the order statistics of the dataset. The AI measures the symmetry in a population by comparing gaps in the data on the left and right sides of the sample median. The TI compares the spread of the dataset at 90% level to the spread at the 75% level. The kurtosis, skewness and the TI reject a Gaussian distribution with a confidence 48 5. Multiple merging in Abell1367 Test Value Rejection of a gaussian AI TI Skewness Kurtosis W -0.077 1.240 0.269 2.680 0.963 ≤ 80 % >99 % >99 % >99 % 98.7 % Table 5.2: 1D substructure indicators for the whole cluster sample level of ≥99%, suggesting that the cluster velocity distribution has longer tails than a Gaussian of the same dispersion. Moreover, in order to assess the normality of the velocity distribution, we use the Wilk - Shapiro (W) test (Yahil & Vidal 1977). Contrary to the χ2 and Kolmogorov Smirnov, this test does not require any hypothesis on the mean and variance of the normal distribution. The W test rejects normality with a confidence level of 98.7%, in agreement with kurtosis, skewness and TI (see Table 5.3). The departure from a normal distribution could result from a mixture of several velocity distributions with different location and smaller velocity dispersion than the whole sample; thus, using the program ROSTAT (Beers et al. 1990), we investigate the presence of significant gaps (Beers et al. 1991) in the velocity distribution, indicating subclustering. A weighted gap is defined by: yi = i(N − i) ∗ (xi+1 − xi ) 1/2 where N is the number of values in the dataset. A weighted gap is significant if its value, relative to the midmean (the mean of the central 50% of the dataset) of the other weighted gaps, is greater than 2.25. This value corresponds to a probability of occurrence in a normal distribution of less than 3%. We detected six significant weighted gaps in the Abell 1367 velocity distribution. The stripe density plot of radial velocities and the position of each gap (indicated with an arrow) are shown in Fig. 5.2. The velocity of the object preceding each gap, the normalized size of the gap and the probability of finding a normalized gap of the same size and position in a normal distribution are listed in Table 5.3. 5.4 Localized velocity structures Given the non-Gaussian nature of the velocity distribution, we looked for spatially localized variations in the LOS velocity and velocity dispersion distributions. First of all we applied the three 3D tests commonly used to quantify the amount of substruc- 5.4. Localized velocity structures 49 Velocity km s−1 Gap Significance 5742 5835 6619 6880 7059 7542 2.53 2.66 2.90 2.64 3.01 2.33 1.40% 1.40% 0.60% 1.40% 0.20% 3.00% Table 5.3: The most significant weighted gaps detected in the velocity distribution of the whole cluster sample. tures in galaxy clusters: the ∆ test (Dressler & Shectman 1988), the α test (West & Bothun 1990) and the test (Bird 1994). The ∆ test is based on the comparison of the local mean velocity, Vlocal , and the velocity dispersion, σlocal , associated to each cluster member (computed using its 10 nearest neighbors) with the mean velocity V , and dispersion σ, of the whole galaxy sample. For each galaxy, the deviation is defined by: δ2 = 11 [(Vlocal − V )2 + (σlocal − σ)2 ] σ2 The observed cumulative deviation ∆, defined as the sum of the δ’s for the cluster members, is used to quantify the presence of substructures. As shown by Pinkney et al. (1996) for samples with no substructures, the value of ∆ is approximately equal to the total number of galaxies, while it is larger in the presence of substructures. The α test measures how much the centroid of the galaxy distribution shifts as a result of correlations between the local kinematics and the projected galaxy distribution. The centroid of the whole galaxy distribution is defined as: xc = N 1 X xi N i=1 yc = N 1 X yi N i=1 For each galaxy i and its 10 nearest neighbors in the velocity space, the spatial centroid is defined as: P11 P11 j=1 xj /σj j=1 yj /σj i i xc = P11 yc = P11 j=1 1/σj j=1 1/σj 50 5. Multiple merging in Abell1367 Indicator Value Prob. of substructures ∆ α 206.5 0.161 Mpc 5.44 1013 M 99.8 % 55.7 % 68.4 % Table 5.4: 3D substructure indicators for our sample where σj is the velocity dispersion for galaxy j and its 10 nearest neighbors in projection. Finally the presence of substructures in the cluster sample is quantified using the α statistic defined as: N 1 X i [(x − xc )2 + (yci − yc )2 ]1/2 α= N i=1 c which represents the mean centroid shift for the galaxy cluster. The higher the value of α, the higher the probability of substructures. The test quantifies the correlations between the position and the projected mass estimator (Heisler et al. 1985), defined as: MP M E = N 32 X v 2 rj πGN j=1 zj where vzj is the radial peculiar velocity with respect to the nearest neighbors group (composed by a galaxy and its 10 nearest neighbors) and rj is the projected distance from the center of the nearest neighbor group. The substructure statistic is then defined as: N 1 X MP M E = Ngal i=1 which represents the average mass of the nearest neighbors groups in the cluster. Since galaxies in the nearest neighbors groups have small projected separations, is generally smaller than the global mass estimate. is lower for a cluster with substructures than for a relaxed system. The value and the significance of the above tests are listed in Table 5.4. These statistical tests are calibrated using 1000 Monte Carlo simulations that randomly shuffle the velocity of galaxies, keeping fixed their observed position, thereby destroying any existing correlation between velocity and position. The probability of subclustering is then given as the fraction of simulated clusters for which the test value is lower (larger for the test) than the observed one. Assuming that these tests reject the null 5.5. The cluster dynamics 51 Figure 5.3: Local deviations from the global kinematics for galaxies in Abell 1367 as measured by the Dressler & Shectman (1988) test. Galaxies are marked with open circles whose radius scales with their local deviation δ from the global kinematics. The ROSAT X-ray contours are shown with dotted lines. hypothesis if the confidence level is greater than 90%, only the ∆ test finds evidence of substructures (see Table 5.4). The local deviations from the global kinematics as measured by the ∆ test are shown in Fig 5.3. The positions of galaxies are marked with open circles whose radius scales with their local deviation δ from the global kinematics. The presence of a substructure with a high deviation from the global cluster kinematic is evident projected near the cluster core. More insights on the cluster dynamical state can be achieved by comparing the results of the one and three dimensional statistical tests with the N-body simulations performed by Pinkney et al. (1996). These authors analyzed how the significance level of statistical tests of substructure varies in different cluster merging scenarios. The deviation of the velocity distribution from a Gaussian and the detection of substructure provided by the ∆ test suggest that Abell 1367 is in the early merging stage, ∼ 0.2 Gyr before core crossing. 5.5 The cluster dynamics The analysis of the galaxy distribution, of the local mean LOS velocity and of the velocity dispersion give further insight onto the cluster structure. The iso-density map of the cluster members (computed using the 10 nearest neighbors to each point) 52 5. Multiple merging in Abell1367 Figure 5.4: Palomar DSS image of the central region (∼1.3 square degrees) of Abell 1367 studied in this Chapter. The iso-density contours for the 146 confirmed cluster members are superposed. The lowest iso-density contour correspond to 3σ above the mean density in the field (left). The ROSAT X-ray contours are superposed in red (right). The straight line indicates the position of the abrupt gas temperature gradient detected by ASCA (Donnelly et al. 1998), used to divide our sample into two subclusters: the North-West and the South-East. is shown in Fig.5.4 (left). The galaxy distribution appears elongated from north-west to south-east with two major density peaks. The highest density region corresponds approximately to the center of the NW X-ray substructure detected by ROSAT (Donnelly et al. 1998), while the secondary density peak is slightly offset from the X-ray cluster center (α(J.2000) = 11h44.8m δ(J.2000) = 19d42m, Donnelly et al. 1998). Moreover the south galaxy density peak roughly coincides with the substructure detected by the ∆ test (see Fig.5.3) and with the infalling group of star-forming galaxies studied by Gavazzi et al. (2003b). The iso-density contours superposed on the ROSAT X-ray contours are shown in Fig.5.4 (right). The region between the two major density peaks coincides with the strong gradient in the gas temperature (see the straight line in Fig.5.4, right) observed for the first time by ASCA (Donnelly et al. 1998) and recently confirmed by Chandra (Sun & Murray 2002). This abrupt temperature change is strongly suggestive of a shock which has generated during a collision between two substructures, probably associated with the SE and the NW galaxy density peaks. In fact N-body simulations show that temperature structures and X-ray morphology similar to the one observed in Abell 1367 are typical of clusters at an early merging phase (∼ 0.25 Gyr before core crossing) (Schindler & Mueller 1993; Gomez et al. 2002). The merging scenario is further supported by the LOS velocity and velocity dispersion fields (computed using the 10 nearest neighbors to each point) shown in Fig. 5.5. The cluster dynamics 53 Figure 5.5: The LOS velocity field (left) and the velocity dispersion field (right) for the whole region studied in this Chapter. The LOS velocity and the velocity dispersion are computed using the 10 nearest neighbors to each pixel, whose size is 36 arcsec2 . The iso-density contours for the 146 confirmed cluster members are superposed in black. 5.5. The SE subcluster has higher LOS velocity and velocity dispersion than the NW substructure. The region with the highest LOS velocity and velocity dispersion lies ∼ 6 arcmin N from the X-ray cluster center and it coincides with the substructure detected by the ∆ test. This result points out the presence of a group of galaxies infalling in the SE cluster core (see Sec.5.5.2). Thus the NW subcluster appears as a relaxed system with the lowest velocity dispersion among the whole sample; on the other hand the SE subcluster appears far from relaxation, and it is probably experiencing a multiple merging event. We use the position of the gas temperature gradient, shown by the straight line in Fig.5.4 (right), to divide our sample into two regions and to study separately the dynamical properties of the two subclusters. A sketch of the cluster dynamical model discussed in the next section is given in Fig.5.6. 5.5.1 The North-West subcluster The NW subcluster is composed of 86 galaxies and includes two density peaks: the highest and a secondary one located at the western periphery of the subcluster (labeled as W subcluster in Fig.5.6), with a weak X-ray counterpart. It has a similar mean location (CBI = 6480 ± 87 km s−1 ) and a lower scale (SBI = 770 ± 60 km s−1 ) than the whole cluster. 54 5. Multiple merging in Abell1367 Figure 5.6: A 3D sketch of Abell 1367 summarizing the various sub-components described in Section 5.5. The cluster is viewed from its near side, as suggested by the eyeball indicating the observer’s position. Fig.5.8 shows the LOS velocity distribution of this subcluster. The W test rejects the Gaussian hypothesis at a confidence level of 39%. Thus the LOS velocity distribution is consistent with a Gaussian distribution, suggesting that this subcluster is a virialized system. Moreover its increasing velocity dispersion profile (see Fig. 5.9) is consistent with a relaxed cluster undergoing two body relaxation in the dense central region, with circular velocities in the center and more isotropic velocities in the external regions (Girardi et al. 1998). However this subcluster also shows some evidences of merging (see Fig.5.7). The brightest galaxy of this cloud CGCG97-095 (NGC3842), located ∼2 arcmin SE from the NW density peak, is a radio galaxy classified as a narrow-angle tail (NAT) (Bliton et al. 1998). The tail orientation (indicated with an arrow in Fig. 5.7) suggests that this galaxy (and the associated substructure) is moving from north-west to southeast, toward the main cluster core. Moreover two CGCG (Zwicky et al. 1961) galaxies, 97-073 and 97-079, show fea- 5.5. The cluster dynamics 55 Figure 5.7: Blow-up of the NW substructure of Abell 1367. The arrows indicate the direction of radio head tails associated with 97-079 and 97-073 and the orientation of the NAT radio galaxy 97-095. The dashed region shows the distribution of the diffuse cluster radio relic (Gavazzi 1978). The iso-density contours for the 146 confirmed cluster members are superposed. tures consistent with the infall scenario. Gavazzi et al. (1995, 2001a) found that both galaxies have their present star formation enhanced along peripheral HII regions which developed at the side facing the direction of motion through the cluster IGM. Their neutral hydrogen is significantly displaced in the opposite side (Dickey & Gavazzi 1991), where 50 kpc long tails are detected both in the light of the synchrotron radiation (Gavazzi & Jaffe 1987) and in Hα (Gavazzi et al. 2001a). The observational scenario is consistent with the idea that ram-pressure (Gunn & Gott 1972) is enhancing for a limited amount of time the star formation of galaxies that are entering the cluster medium for the first time. However these two galaxies appear not directly associated with the center of the NW subcluster since they lie at a projected distance of ∼0.34 Mpc from the main density peak (see Fig.5.7). Moreover their large distance (∼0.48 Mpc) from the shock front observed in X-ray between the NW and the SE substructure indicates that these objects do not belong to the main galaxy density peak infalling into the cluster center. Conversely they are at a projected distance of only 0.08 Mpc from the center of the W subcluster, suggesting that they are associated with this subcloud. For these reasons we consider an alternative scenario in which these two galaxies belong to a secondary substructure infalling into the NW substructure from the western side (see Fig. 5.6). This picture is supported by the presence of the extended radio relic detected both in X-ray and radio continuum in this region (Gavazzi 1978; Gavazzi & Trinchieri 1983). Cluster radio halos contain fossil radio plasma, the former outflow of a radio galaxy, that has been revived by shock compression during cluster merging 56 5. Multiple merging in Abell1367 Figure 5.8: The LOS velocity distribution for galaxies in the NW (upper) and in the SE (lower) subclusters. (Enßlin et al. 1998; Enßlin & Brüggen 2002). The radio relic observed in Abell 1367 extends, south-west to north-east, from 97-073 to 127-040 with a projected extent of 0.8 Mpc (see Fig.5.7). The age of its electrons is estimated to be ∼ 0.2 Gyr (Enßlin et al. 1998). The only plausible source of high energy electrons available in this region is the NAT galaxy 97-095, presently at ∼0.25 Mpc from the relic and whose tails point exactly in the relic direction. Assuming that the fossil radio halo originated from 97-095, we find that the infall velocity of this galaxy into the SE subcluster is V ∼ 1250 km s−1 , consistent with the typical infall velocity of cluster galaxies. Thus the presence of the radio relic results consistent with a merging scenario in which the W subcluster, containing 97-079 and 97-073, is infalling into the NW substructure, compressing the plasma ejected from 97-095 and re-accelerating the electrons to relativistic energies. 5.5.2 The South-East subcluster The SE cloud is composed of 60 galaxies associated with the X-ray cluster center. It has the highest LOS velocity and dispersion of the whole sample (see Fig.5.5) with a location CBI = 6596 ± 137 km s−1 and a scale SBI = 1001 ± 70 km s−1 . Its velocity distribution, shown in Fig. 5.8, appears significantly non-Gaussian. The W test rejects the Gaussian hypothesis at a confidence level of 96.8%, supporting the idea that 5.5. The cluster dynamics 57 Figure 5.9: The velocity dispersion radial profile of the NW (upper) and the SE (lower) subclusters. the cluster center is far from relaxation. This is in agreement with the decreasing velocity dispersion profile of this region (see Fig.5.9), consistent with isotropic velocities in the center and radial velocities in the external regions, as expected in the case of galaxy infall onto the cluster (Girardi et al. 1998). The velocity distribution of Fig. 5.8 has three peaks at ∼ 5500 km s−1 , ∼ 6500 km s−1 and ∼ 8200 km s−1 respectively, probably associated with three separate groups. Moreover we remark that the galaxy gaps between the three peaks are fairly consistent with two of the most significant weighted gaps detected in the global velocity distribution (V ∼ 5800 km s−1 and V ∼ 7500 km s−1 ). In order to check for any position-velocity segregation, we divide the SE subcluster in three groups according to their LOS velocity: galaxies with V < 5800 km s−1 belong to the low velocity group, galaxies with V > 7500 km s−1 belong to the high velocity group and galaxies with intermediate velocity belong to the SE subcluster. The projected distribution of the three groups is shown in Fig.5.10. The high-velocity group (V ∼ 8200 km s−1 , triangles) appears segregated in the northern part of the SE cloud, extending ∼20 arcmin in right ascension but only ∼7 arcmin in declination. It is associated with the substructure detected by the ∆ test (see Fig. 5.3) and with the infalling group of star-forming galaxies recently discovered by Sakai et al. (2002) and by Gavazzi et al. (2003b). Its spatial segregation and high star formation activity suggest that this group is a separate unit infalling into the cluster, probably from the 58 5. Multiple merging in Abell1367 Figure 5.10: The distribution of galaxies belonging to the South-East subcluster. Triangles indicate galaxies with LOS velocity > 7500 km s−1 , circles galaxies with LOS velocity < 5800 km s−1 and squares galaxies with LOS velocity comprises in the range 5800 km s−1 < V < 7500 km s−1 . The ROSAT X-ray contours are shown. near side (see Fig. 5.6). It is remarkable that Sun & Murray (2002), using Chandra observations of the cluster center, discovered a ridge-like structure around the cluster center, ∼6 arcmin south from the center of the high velocity group, probably associated with a compact merging subcluster (perhaps this group) penetrating the SE cluster core. The low-velocity group (V ∼ 5500 km s−1 , circles in Fig.5.10) seems segregated in the eastern part of the cloud, perhaps infalling from the eastern side into the cluster core (Fig. 5.6). This scenario is also supported by the detection of cool gas streaming into the cluster core from the eastern side (Forman et al. 2003), probably associated with this low velocity group of galaxies. Galaxies with V ∼ 6500 km s−1 (squares in Fig.5.10) are homogeneously distributed over the SE subcluster, representing its virialized galaxy population. However the brightest galaxy in this group 97-127 (NGC3862) is a NAT radio galaxy with very extended radio tails pointing in the direction of the low velocity group (Gavazzi et al. 1981), suggesting motion relative to the IGM. The velocity-space segregation observed in the SE subcluster suggests that the cluster center is experiencing multiple merging of at least two separate groups, supporting the idea that it is far from relaxation. This picture is consistent with the high gas 5.6. Star formation activity in the infalling groups 59 Figure 5.11: The LOS velocity distribution for emission line (upper) and non emission line galaxies (lower) in the whole cluster sample. entropy in this region, since in absence of a cool dense core the substructures infalling into the major cluster can penetrate deep inside, disturbing the cluster core dynamics (Churazov et al. 2003). A sketch of the various substructures identified in Abell 1367 by the present study, is given in Fig. 5.6. Five substructures are detected. Two clouds, the NW and SE subclusters, are in the early merging phase, meanwhile three smaller groups are infalling into Abell 1367. The W subcloud, associated with the head-tail systems 97-073/79, is probably infalling into the NW subcluster, exciting the radio relic observed in between the two structures. The other two groups are infalling into the SE subcluster: the low velocity group from the eastern side, while the high velocity group from the near side. 5.6 Star formation activity in the infalling groups The dynamical study presented in the previous sections indicates that Abell 1367 is a dynamically young cluster in the early stage of a multiple merging event involving at least five substructures. Since merging is expected to trigger star formation in cluster galaxies (Bekki 1999), we study separately the spatial and velocity distribution of the star forming galaxies. Only 49 out of the 146 cluster members show recent star 60 5. Multiple merging in Abell1367 Figure 5.12: Projected density map of non emission line (left) and emission line (right) galaxies in Abell 1367. The iso-density contours of the 146 confirmed cluster members are superposed. formation activity (e.g. Hα line in emission, Iglesias-Páramo et al. 2002; Gavazzi et al. 2003a; Cortese et al., in preparation). Fig.5.11 shows the LOS velocity distribution of galaxies divided into emission line (upper panel) and non emission line (lower panel) galaxies. The star forming sample has higher location and scale (CBI = 6704 ± 168 km s−1 , SBI = 1076 ± 76 km s−1 ) than the quiescent sample (CBI = 6446 ± 79 km s−1 , SBI = 738 ± 58 km s−1 ). According to a two-sample Kolmogorov-Smirnov test the two velocity distributions have only ∼5% probability of being consistent, suggesting a different origin and/or evolution. We remark that, if the star forming galaxies are infalling onto the cluster along radial orbits, their velocity dispersion √ should be ∼ 2 times the velocity dispersion of the relaxed sample, as observed in this case. This result suggests that star forming systems are an infalling population while the non-star forming galaxies represent the virialized cluster population. The projected density distribution of star forming and non star forming is shown in Fig.5.12. The highest density of non emission line systems is observed near the center of the NW substructure. This morphological segregation further supports the idea that the NW cloud is a relaxed system merging for the first time into the SE subcluster. The emission line galaxies have a different distribution. The highest density of star forming systems is in the infalling groups, i.e. in the high velocity group infalling into the SE subcluster and in the W cloud infalling into the NW substructure, suggesting that their interaction with the cluster environment is triggering some star formation activity. Indeed in these systems the fraction of star forming galaxies lies between 64% and 36%, decreasing to 31% in the NW substructure and to 20% in the SE subcluster. 5.7. Cluster mass 61 Sample RH Mpc A1367 A1367 A1367 A1367 A1367 A1367 0.41 0.37 0.30 0.24 0.27 0.26 all types non-star forming NW all types NW non-star forming SE all types SE non-star forming MV 1014 M 7.04 4.35 3.87 2.47 5.80 3.90 ± ± ± ± ± ± 0.90 0.70 0.62 0.46 0.88 0.83 MPM 1014 M 7.82 5.11 6.12 3.29 6.87 5.58 ± ± ± ± ± ± 2.50 0.90 1.52 0.59 1.20 0.74 Table 5.5: Mass estimate for Abell 1367 5.7 Cluster mass The virial theorem is the standard tool used to estimate the dynamical mass of galaxy clusters. Under the assumptions of spherical symmetry and hydrostatic equilibrium and if the mass distribution follows the distribution of the observed galaxies independent of their luminosity, the total gravitational mass of a cluster is given by MV = 3π 2 σ RH G where σ is the galaxy velocity dispersion and RH is the cluster mean harmonic radius: N (N − 1) RH = P −1 i>j Rij where N is the total number of galaxies. An alternative approach is to use the projected mass estimator (Heisler et al. 1985), defined as 32 X 2 MP M = V Ri πGN i i where Vi is the observed radial component of the velocity of the i galaxy with respect to the systemic cluster velocity, and Ri is its projected separation from the cluster center. The numerical factor 32 assumes that galaxy orbits are isotropic. In case of purely radial or purely circular orbits this factor becomes 64 or 16 respectively. Mass estimates obtained using the two above methods and their uncertainties are listed in Table 5.7. We remark that these mass estimates are probably biased by the dynamical state of Abell 1367, which appears far from virialization. In particular the presence of substructures leads to an overestimate of the cluster mean harmonic radius and velocity dispersion, and thus of the virial mass (Pinkney et al. 1996). For 62 5. Multiple merging in Abell1367 this reason the mass derived for the whole cluster and for the SE and NW subclusters separately is probably overestimated. Assuming that the early type sample represents the virialized cluster population (see previous section), we also derive mass estimates for the three dynamical units using the non star forming systems only. For all the studied samples the virial mass estimates are affected by smaller uncertainties and yield smaller values than the projected mass estimates. This can be due to the contamination by interlopers (Heisler et al. 1985) or, more probably, to the assumption of isotropic orbits. Indeed assuming purely radial or circular orbits the mass estimate varies by a factor of 2, becoming consistent with the virial mass. The mass inferred from the non-star forming population are, as expected, systematically lower than the ones obtained from all types. The value obtained for the whole sample is consistent with the mass estimates available in the literature (MV = 7.26 ± 1.40 1014 M Girardi et al. 1998; MV = 6.07±0.93 1014 M , MP M = 6.28±0.80 1014 M Rines et al. 2003). 5.8 Two-Body Analysis In this section we investigate whether the two clouds A1367NW, A1367SE and the three groups infalling into the SE and NW subclusters form gravitationally bound systems. For each system we apply the two-body analysis described by Beers et al. (1991). The two subclumps are treated as point masses moving on radial orbits. They are assumed to start their evolution at time t=0 with zero separation, and are moving apart or coming together for the first time in their history. For bound radial orbits, the parametric solutions to the equations of motion are: Rm (1 − cos χ) 2 R3 1/2 m t= (χ − sin χ) 8GM 2GM 1/2 sin χ V = Rm (1 − cos χ) R= where R is the components separation at time t, and V is their relative velocity. Rm is the separation of the subclusters at maximum expansion and M is the total mass of the system. Similarly, the parametric solutions for the unbound case are: R= GM (cosh χ − 1) 2 V∞ t= GM (sinh χ − χ) V∞3 5.8. Two-Body Analysis 63 Figure 5.13: The bound and unbound orbit regions in the (Vrel , α) plane. The bound-incoming solutions (BIa and BIb ), the bound-outgoing solutions (BO) and the unbound-outgoing (UO) solutions are indicated with solid lines. The dotted lines show the dividing line between bound and unbound regions. The vertical solid lines represent the observed Vrel and the dashed regions their associated 1σ uncertainty. 64 5. Multiple merging in Abell1367 V = V∞ sinh χ (cosh χ − 1) where V∞ is the asymptotic expansion velocity. The system parameters V and R are related to the observables Vrel (the LOS relative velocity) and Rp (the projected separation) by: Vrel = V sin α, Rp = R cos α where α is the angle between the plane of the sky and the line joining the centers of the two components. The two systems are thus closed by setting the present time to t0 = 13 Gyr (the age of the Universe in a Ωm =0.3 and Ωλ =0.7 cosmology) and solved iteratively to determine the projection angle as a function of Vrel . We determine two solutions for each two-body model, assuming two extreme values for the total mass of each system ranging from the virial mass of the non-star forming population to the virial mass of the whole cluster. Table 5.8 summarizes the adopted parameters of the two-body analysis, and Fig. 5.13 shows the computed solutions in the (α, Vrel ) plane. The vertical lines represent the observed values of Vrel and the dashed regions their associated 1σ uncertainties. The solutions have three different regimes: an unbound-outgoing regime (UO), a bound-outgoing regime (BO) and a bound-ingoing regime (BI). It is easy to show that the unbound solutions will lie in the region of the (α, Vrel ) plane where: 2 Vrel Rp ≤ 2GMtot sin2 α cos α. The dotted lines in Fig. 5.13 show the dividing line between bound and unbound regions. In the BO regime, the two subclumps are still separating and have not yet reached the maximum expansion. The BI regime describes the system after maximum expansion. For each Vrel , there are two corresponding values of α, a large and a small one. The large value assumes that the substructures are far apart, with low relative velocity, while the small value implies that the subclusters are close together near the plane of the sky (see Fig. 7 in Beers et al. 1991). Thus we split the BI regime into two branches, called BIa and BIb . The probability of each solution, computed following the procedure described by Beers et al. (1991), is given in Table 5.8. Our result is that the A1367NW/SE and the A1367SE/High Velocity group systems are bound with 100% probability and presently infalling with 96% and 100% probability respectively. The A1367NW/W and the A1367SE/Low Velocity group systems are bound at 99% and 96% probability respectively. We conclude that all systems constituting Abell 1367 are gravitationally bound at ≥ 96% probability. 5.9. Conclusions 65 System A1367NW/SE A1367NW/W A1367SE/Low Vel. gr. A1367SE/High Vel. gr. Mtot Vrel ± ∆Vrel Rp 1014 M km s−1 Mpc Solution Probability BIa BIb BO UO % % % % 7.04 4.35 7.04 2.47 7.04 3.90 7.04 3.90 84 ± 162 84 ± 162 500 ± 200 500 ± 200 1000 ± 200 1000 ± 200 1500 ± 200 1500 ± 200 0.45 0.45 0.37 0.37 0.38 0.38 0.08 0.08 57 55 57 56 58 57 56 58 40 41 40 41 40 39 44 42 3 4 2 2 0 0 0 0 0 0 1 1 2 4 0 0 Table 5.6: Two-body model parameters 5.9 Conclusions I have presented a dynamical analysis of the central ∼ 1.3 square degrees of the galaxy cluster Abell 1367, based on 273 redshift of which 119 are new measurements. The LOS velocity distribution of the 146 cluster members is significantly non Gaussian, suggesting that the cluster is dynamically young. The member galaxies show an elongated distribution along the NW-SE direction with two major density peaks, consistent with the X-ray morphology. The strong difference in the LOS velocity and velocity dispersion of the two density peaks, the abrupt gas temperature gradient detected in X-rays and the 3D statistical tests support a merging scenario involving at least two subclusters. Moreover the dynamical properties of the NW and SE clouds suggest an even more complex picture, summarized in Fig. 5.6. At least another group of star forming galaxies (the high velocity group) infalling into the cluster core is detected, suggesting a multiple merging event. Furthermore our analysis suggests the presence of two other groups infalling into the cluster center. In the North-West part of Abell 1367 a group of galaxies (W subcluster), associated with the infalling galaxies 97-073/79 and with the radio relic observed in this region, is probably merging with the relaxed core of the NW subcluster. In the South part another group (the low velocity group) is infalling from the eastern side into the disturbed core of the SE subcluster. These three subgroups have a higher fraction of star forming galaxies than the cluster core, as expected during the early phase of merging events. The multiple merging scenario is consistent with the location of Abell 1367 being at the intersection of two filaments, the first extending roughly 100 Mpc from Abell 1367 toward Virgo (West & Blakeslee 2000) and the second extending between Abell 1367 66 5. Multiple merging in Abell1367 and Coma (as a part of the Great Wall, Zabludoff et al. 1993). As predicted by Katz & White (1993) this is the natural place for Abell 1367 to evolve into a rich relaxed cluster. 5.9. Conclusions 67 Name 114000+195426 114159+193227 114200+195846 114208+191905 114212+195650 114213+193001 114215+200427 114219+200548 114224+195329 114224+191157 114226+194317 114230+191447 114230+192553 114238+194718 114239+195145 114240+195627 114243+191615 114249+193935 114250+193955 114252+195656 114254+193851 114254+194033 114258+194321 114258+194053 114258+194644 114258+195612 114259+194801 114300+192515 114301+194758 114301+195313 114307+192807 114307+193029 114310+192526 114310+191519 114313+200747 114314+194821 114314+192534 114317+195525 114317+194658 114318+201523 114319+192520 114320+193637 114320+195206 114322+195704 114324+194121 114332+201326 114332+195108 114335+200005 114336+193930 114337+193835 114337+201533 114339+193446 114342+193636 114343+195607 114345+201252 114350+195702 114350+194138 114353+195004 114353+194422 114353+194315 R.A. (J.2000) Dec. (J.2000) r’ mag V km s−1 Tel. 114000.62 114159.52 114200.83 114208.01 114212.47 114213.87 114215.59 114219.15 114224.39 114224.48 114226.24 114230.62 114230.95 114238.24 114239.78 114240.26 114243.81 114249.85 114250.47 114252.17 114254.40 114254.93 114258.13 114258.37 114258.53 114258.94 114259.71 114300.65 114301.24 114301.97 114307.13 114307.16 114310.09 114310.29 114313.18 114314.49 114314.99 114317.25 114317.61 114318.05 114319.68 114320.44 114320.66 114322.06 114324.66 114332.24 114332.72 114335.47 114336.07 114337.17 114337.82 114339.09 114342.18 114343.12 114345.50 114350.16 114350.83 114353.42 114353.45 114353.61 195426.7 193227.3 195846.0 191905.0 195650.3 193001.6 200427.0 200548.0 195329.8 191157.0 194317.1 191447.5 192553.8 194718.6 195145.9 195627.5 191615.8 193935.1 193955.7 195656.4 193851.3 194033.6 194321.1 194053.9 194644.2 195612.7 194801.1 192515.2 194758.9 195313.5 192807.3 193029.8 192526.4 191519.2 200747.9 194821.7 192534.3 195525.1 194658.2 201523.3 192520.9 193637.1 195206.2 195704.7 194121.4 201326.1 195108.2 200005.6 193930.8 193835.8 201533.5 193446.2 193636.3 195607.8 201252.2 195702.0 194138.0 195004.6 194422.2 194315.8 15.98 15.63 17.09 19.04 17.73 16.92 19.20 16.45 18.29 16.39 17.50 18.60 17.80 17.04 19.06 18.63 18.72 19.14 19.22 16.69 17.17 18.87 18.98 19.25 19.00 18.41 18.92 18.42 18.67 18.61 17.37 17.93 16.62 17.41 16.40 19.29 15.76 18.88 15.69 17.81 15.99 19.71 18.15 16.94 18.33 16.46 19.07 16.38 19.24 17.31 20.19 16.01 19.26 18.56 19.27 17.98 19.13 19.23 15.66 17.17 10883 21228 6420 23456 20278 23641 6100 6841 31440 28546 23416 27304 45683 25610 53710 19946 5312 72429 13759 5936 6406 71389 6523 71436 88274 7059 71600 53145 72572 46935 32298 23763 19188 23578 5383 71433 23867 30273 6295 46170 6757 44171 52416 7909 35778 33438 14313 20600 44616 12502 11464 7477 71296 19711 20476 6848 72744 27946 6141 23578 CAN CAN WHT WHT WHT WHT WHT CAN WHT CAN WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT CAN WHT CAN WHT CAN WHT WHT WHT WHT WHT WHT CAN WHT CAN WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT Table 5.7: The 119 new redshift measurements 68 5. Multiple merging in Abell1367 Name 114356+201404 114357+201122 114358+195330 114359+195630 114402+194742 114403+200552 114404+192922 114404+195956 114407+193850 114407+193143 114412+195503 114412+195633 114412+201119 114415+193037 114415+193012 114417+194543 114422+194628 114426+195951 114430+194258 114432+195341 114432+194734 114447+201248 114449+195628 114501+195504 114503+193831 114503+194743 114504+201412 114505+194057 114506+200849 114509+194845 114509+193316 114509+194526 114516+193245 114517+200120 114517+201108 114517+200110 114520+194220 114520+193259 114522+195146 114524+201239 114526+201056 114529+195658 114530+193639 114531+200217 114533+194505 114533+200028 114536+194253 114540+194302 114543+193854 114543+193905 114544+194013 114545+193151 114545+201200 114548+192708 114549+195915 114550+194824 114602+194754 114605+195151 114620+194518 R.A. (J.2000) Dec. (J.2000) r’ mag V km s−1 Tel. 114356.80 114357.69 114358.86 114359.51 114402.65 114403.70 114404.17 114404.65 114407.21 114407.71 114412.22 114412.27 114412.92 114415.25 114415.33 114417.28 114422.16 114426.10 114430.30 114432.19 114432.98 114447.20 114449.72 114501.97 114503.00 114503.14 114504.25 114504.83 114506.38 114509.38 114509.40 114509.65 114516.18 114517.10 114517.29 114517.64 114520.33 114520.49 114522.62 114524.33 114526.27 114529.39 114530.37 114531.31 114533.88 114533.97 114536.19 114540.32 114543.65 114543.77 114544.86 114545.66 114545.78 114548.13 114549.88 114550.61 114602.12 114605.35 114620.85 201404.9 201122.7 195330.2 195630.8 194742.7 200552.6 192922.8 195956.6 193850.9 193143.1 195503.9 195633.4 201119.7 193037.5 193012.3 194543.9 194628.2 195951.5 194258.3 195341.6 194734.6 201248.5 195628.9 195504.5 193831.2 194743.9 201412.2 194056.9 200849.9 194845.4 193316.2 194526.9 193245.1 200120.7 201108.8 200110.0 194220.3 193259.4 195146.5 201239.3 201056.8 195658.2 193639.4 200217.5 194505.9 200028.7 194253.7 194302.8 193854.9 193905.9 194013.3 193151.4 201200.3 192708.4 195915.3 194824.6 194754.3 195151.0 194518.0 18.42 17.06 19.22 20.37 17.52 15.80 18.59 17.33 17.10 18.44 17.65 17.02 19.25 16.58 18.27 18.14 15.70 16.98 18.78 18.89 18.82 18.17 16.70 18.79 16.76 17.91 18.31 15.67 19.23 17.49 15.80 16.90 16.80 15.32 18.21 15.46 20.44 17.48 21.14 18.73 16.40 16.29 17.20 19.78 18.11 17.56 18.60 17.74 16.93 16.30 17.18 18.57 19.11 16.72 15.87 18.85 19.62 18.86 17.58 72058 5348 6200 6992 43665 5698 53335 33830 20877 53424 20916 6244 74731 6502 35227 66264 6527 30102 40347 42649 71100 6699 5539 45708 6193 23374 5477 6506 3822 19831 7409 19834 19669 14745 79253 14713 54544 4653 18012 44376 20134 24000 40000 45691 31440 35830 48966 5545 7828 7301 19487 6880 27431 30193 20035 41484 73746 46635 45683 WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT WHT CAN WHT WHT WHT WHT WHT WHT WHT LOI WHT WHT LOI WHT WHT LOI WHT WHT LOI WHT LOI WHT WHT WHT WHT CAN WHT LOI WHT WHT WHT WHT WHT LOI LOI WHT LOI WHT WHT CAN WHT WHT WHT WHT Table 5.7: Continue Chapter 6 Unveiling the evolution of early type galaxies with GALEX. 6.1 Introduction In Chapters 3 and 4 I have shown that at low UV luminosities the contribution of early-type quiescent galaxies is not negligible. This represents the first evidence of a morphology/star formation - density relation at ultraviolet wavelengths and demonstrates that we cannot blindly assume all UV selected galaxies are star-forming systems, especially at low UV luminosities and in high density environments. This also points out the strong potential of ultraviolet observations for studying all cluster galaxies: not only star-forming systems in which UV emission traces the presence of newly born stars, but also early type galaxies whose emission is usually ascribed to low mass old post asymptotic giant branch stars. The excess ultraviolet radiation from giant early-type galaxies is in fact supposed to arise from hot low mass stars in late stages of stellar evolution (O’Connell 1999). All theoretical, spectral and imaging evidences have recently converged towards the view that the UV emission originates from He-burning, extreme horizontal branch stars, their post-HB progeny and post-AGB stars in the dominant, metal rich stellar population of elliptical galaxies. However it is still unknown whether the UV emission of all early type galaxies is dominated by the contribution of old stellar populations independently from the galaxy morphology (i.e. ellipticals vs. lenticulars) and luminosity (i.e. dEs vs. giant Es). In particular it would be interesting to know if the UV properties of dwarf elliptical galaxies differ from those of giants, as much as other structural (Gavazzi et al. 2005) and kinematic (van Zee et al. 2004) properties depend on luminosity, due to their different star formation histories (single episodic vs. burst) (Ferguson & Binggeli 1994; Grebel 2000). In fact, a recent burst of star formation would strongly contribute to the UV emission of an elliptical galaxy, even if its stellar population is 69 70 6. Unveiling the evolution of early type galaxies with GALEX. dominated by old low mass stars. Due to morphological segregation (Whitmore et al. 1993), nearby clusters are the ideal targets for assembling complete, volume limited samples of early-type objects. As part of a study aimed at analyzing the environmental dependence of galaxy evolution, we observed large portions of the Virgo cluster with GALEX (Boselli et al. 2005a). Owing to the superior quality of the photographic material obtained by Sandage and collaborators, an extremely accurate and homogeneous morphological classification exists for Virgo galaxies, down to mB ≤ 18 mag (MB ≤-13 assuming a distance of 17 Mpc), allowing a detailed discrimination among different subclasses of early-type galaxies (ellipticals, lenticulars, dwarfs) and from quiescent spirals. Furthermore a wealth of ancillary data for many Virgo members, covering a large portion of the electromagnetic spectrum from the visible to the infrared is available from the GOLDMine database (Gavazzi et al. 2003a). 6.2 Data The analysis presented in this Chapter is based on an optically selected sample of early-type galaxies including giant and dwarf systems (E, S0, S0a, dE and dS0) extracted from the Virgo Cluster Catalogue of Binggeli et al. (1985), which is complete to mB ≤18 mag (MB ≤ -13). The Virgo cluster region was observed in spring 2004 as part of the All Imaging Survey (AIS) and of the Nearby Galaxy Survey (NGS) carried out by the Galaxy Evolution Explorer (GALEX) in two UV bands: FUV (λeff = 1528Å, ∆λ = 442Å) and NUV (λeff = 2271Å, ∆λ = 1060Å), covering 427 objects. See Chapter 2 and Martin et al. (2005) and Morrissey et al. (2005) for details on the GALEX instrument and data characteristics. The present sample includes all Virgo cluster early-type systems detected in the NUV GALEX band (264 objects, 194 from the NGS); of these, 126 (of which 74 from the NGS) have been also detected in the FUV. The resulting sample is thus ideal for the proposed analysis as it provides us with the first large volume-limited sample of elliptical, lenticular and dwarf galaxies spanning 4 dex in luminosity with homogeneous data. Whenever available, we extracted fluxes from the deep NGS images, obtained with an average integration time of ∼ 1500 sec, complete to mAB ∼ 21.5 in the NUV and FUV. Elsewhere UV fluxes have been extracted from the less deep AIS images (∼ 70 sq. degrees), obtained with an average integration time of ∼ 100 sec, complete to mAB ∼ 20 in both the FUV and NUV bands. The resulting sample, although not complete in both UV bands, includes giants and dwarf systems: at a limiting magnitude of MB ≤ -15, 71 % of the observed galaxies have been detected in the NUV, 46% in the FUV. All UV images come from the GALEX IR1.0 release. UV fluxes were obtained by integrating GALEX images within elliptical annuli of increasing diameter up to the optical B band 25 mag arcsec−2 isophotal radii consistently with the optical 6.3. The UV properties of early-type galaxies 71 and near-IR images. Independent measurements of the same galaxies obtained in different exposures give consistent photometric results within 10% in the NUV and 15% in the FUV in the AIS, and about a factor of two better for bright (NUV ≤16) galaxies. The statistical uncertainty in the UV photometry is on average a factor of ∼ 2 better in the NGS than in the AIS especially for fainter objects. UV data have been combined with multifrequency data taken from the GOLDMine database (http://goldmine.mib.infn.it; Gavazzi et al. 2003a). These are B and V imaging data, mostly from Gavazzi et al. (2005) and Boselli et al. (2003a), and nearIR H imaging from Gavazzi et al. (2000, 2001c). Optical and near-IR data have on average a photometric precision of ∼ 10%. Spectroscopic metallicity index Mg2 and velocity dispersion data come from GOLDMine or from Golev & Prugniel (1998) and Bernardi et al. (2002). Galaxies analyzed in this Chapter are all bona-fide Virgo cluster members: given the 3-D structure of the cluster, distances have been assigned following the subcluster membership criteria of Gavazzi et al. (1999a). Owing to the high galactic latitude of Virgo, no galactic extinction correction was applied (AB ≤ 0.05). 6.3 The UV properties of early-type galaxies Despite the complex 3-D structure of Virgo (Gavazzi et al. 1999a), the uncertainty on the distance (hence on the luminosity) of the target galaxies, does not constitute a major source of dispersion in the determination of the color-magnitude (CMR) relation. Figure 6.1 shows various UV to optical and near-IR CMRs. Similar results are obtained if, instead of the mass-tracer H band luminosity (Zibetti et al. 2002), we use the B band absolute magnitude. The NUV to optical (Fig. 6.1b) and near-IR (Fig. 6.1a) CMRs are well defined and are similar to optical or near-IR CMRs, with brighter galaxies having redder colors, independent of their morphological type: the color index (N U V − V ) increases by ∼ 2 magnitudes from dwarfs (LH ∼ 108 LH ) to giants (LH ∼ 1011.5 LH ), while (N U V − H) changes by ∼ 3 mag. A weak flattening of the relation appears for LH ≥ 1010 LH . This behavior confirms the one reported by Ferguson (1994) in the (B − V ) vs. MB CMR. On the contrary, the FUV to optical (Fig. 6.1d) and near-IR (Fig. 6.1c) CMRs differ systematically for dwarfs and giant systems: galaxies brighter than LH ∼ 109.5 LH have similar red colors, while for LH ≤ 109.5 LH colors become progressively bluer. Even if this trend can be due to a selection effect, (reddest dwarfs being undetectable in the FUV), it is indisputable that there exists a significant population of dEs with bluer colors than Es and S0s. The dichotomy between giants and dwarfs is even more apparent in the UV color index (F U V − N U V ) (see Fig. 6.2). The (F U V − N U V ) becomes redder with increasing luminosity for dwarf ellipticals while, on the contrary, it becomes bluer for giant ellipticals (Fig. 6.2a). The blueing relation is tight among 72 6. Unveiling the evolution of early type galaxies with GALEX. Figure 6.1: The near-UV (left column) and far-UV (right column) to optical and nearIR color magnitude relations. Colors are in the AB magnitude system. Open circles are for ellipticals, filled circles for dwarfs, crosses for lenticulars (S0-S0a). Galaxies redder than the dashed line are undetectable by the present survey (at the NGS limit). Largest 1σ errors for luminous and dwarf systems are given. 6.3. The UV properties of early-type galaxies 73 Table 6.1: Main relations for early type galaxies x y a b R rms Ellipticals1 LH LH LH LH LH B−H σ F UV − NUV F UV − H NUV − H F UV − V NUV − V F UV − NUV F UV − NUV −0.30 ± 0.14 −0.22 ± 0.19 0.17 ± 0.18 −0.15 ± 0.18 0.26 ± 0.12 −0.84 ± 0.45 −1.35 ± 0.37 Lenticulars +4.52 ± 1.52 +10.55 ± 2.10 +4.85 ± 1.85 +8.38 ± 1.88 +2.55 ± 1.30 +3.22 ± 0.98 +4.39 ± 0.89 −0.47 −0.28 0.22 −0.21 0.45 −0.43 −0.69 0.31 0.43 0.47 0.38 0.31 0.32 0.26 LH LH LH LH LH B−H σ F UV − NUV F UV − H NUV − H F UV − V NUV − V F UV − NUV F UV − NUV −0.28 ± 0.15 0.31 ± 0.21 0.61 ± 0.11 0.03 ± 0.23 0.49 ± 0.09 −1.00 ± 0.32 −1.29 ± 0.39 Dwarf s +4.40 ± 1.62 +0.75 ± 2.00 +0.51 ± 1.17 +6.62 ± 2.38 +0.26 ± 1.00 +3.70 ± 0.70 +4.28 ± 0.84 −0.31 0.27 0.65 0.03 0.68 −0.49 −0.58 0.45 0.58 0.36 0.59 0.25 0.42 0.39 LH LH LH LH LH B−H σ F UV − NUV F UV − H NUV − H F UV − V NUV − V F UV − NUV F UV − NUV 1.73 ± 0.41 2.55∗ ± 0.55 0.91 ± 0.19 1.91∗ ± 0.55 0.63 ± 0.17 0.95 ± 0.45 − −13.90 ± 2.16 −15.97∗ ± 4.96 −2.72 ± 1.68 +11.35∗ ± 4.93 +1.28 ± 1.05 +0.12 ± 0.73 − 0.52 0.68∗ 0.56 0.60∗ 0.49 0.40 − 0.59 0.91∗ 0.57 0.87∗ 0.47 0.60 Notes to Table: Col. 1 and 2: x and y variables Col. 3 and 4: slope a and intercept b of the bisector linear fit with weighted variables Col. 5: Pearson correlation coefficient Col. 6: mean dispersion around the best fit 1: excluding VCC 1499 *: uncertain values because of the UV detection limit 74 6. Unveiling the evolution of early type galaxies with GALEX. Figure 6.2: The relationship between the UV color index (F U V − N U V ) and a) the total H band luminosity, b) the B-H color index, c) the logarithm of the central velocity dispersion and d) the Mg2 index. Symbols are as in fig. 6.1. Labeled points indicate objects having unusual radio or optical properties (see Sect. 3). 6.4. Discussion and conclusion 75 ellipticals (see Table 1) and barely observed in lenticulars because of their higher dispersion 1 . A similar behavior between ellipticals and lenticulars is observed in the (F U V −N U V ) color relation (Fig. 6.2b): this mixed giant population becomes bluer in the UV with increasing reddening in the (B − H) color index. The behavior of dwarf ellipticals is different: although with a huge dispersion, the (F U V −N U V ) color index reddens as the (B −H) and the other optical color indexes. The dichotomy between dwarf and giant systems cannot be observed in the run of (F U V − N U V ) vs. central velocity dispersion (which is directly related to the system total dynamical mass; Fig. 6.2c) nor as a function of the metallicity sensitive (Poggianti et al. 2001) Mg2 Lick index (Fig. 6.2d) because these two parameters are not available for dwarfs. In ellipticals and lenticulars the UV color index (F U V − N U V ) depends on both the metallicity index Mg2 and σ in a way opposite to the behavior at optical wavelengths, where galaxies are redder when having higher Mg2 and velocity dispersions. 6.4 Discussion and conclusion For the first time the UV properties of early-type galaxies have been studied down to MB ∼ -15 mag. The comparison with previous studies is thus limited to the brightest objects. Our CMR can be compared with the one obtained by Yi et al. (2005) based on a complete sample of bright early-type objects (Mr ≤ -20 mag) extracted from the Sloan Digital Sky Survey (SDSS) by Bernardi et al. (2003). The CMR presented by Yi et al. (2005) (N U V − r vs. Mr ) shows a significantly larger dispersion (σ ≥ 1.5 mag) than the one found in Virgo (see Table 1). As discussed in Yi et al. (2005), the large dispersion in their CMR can be ascribed to galaxies with a mild or residual star formation activity included in the Bernardi et al. (2003) sample. If restricted to the ”UV weak” sample, the dispersion in the Yi et al. relation drops to 0.58 mag, i.e. still larger than the one seen in the Virgo cluster in the same luminosity range. Despite possible larger distance uncertainties in the SDSS, the difference in the scatter between our and the Yi et al. (2005) CMR might arise from the classification in the SDSS that uses concentration indices and luminosity profiles in discriminating hot from rotating systems. It is in fact conceivable that the larger dispersion in the CMR of ”UV weak” galaxies of Yi et al. (2005) comes from the contamination of qui1 The scatter in the blueing relation among ellipticals decreases significantly (from 0.31 to 0.10) if we exclude the misclassified post-starburst dwarf VCC 1499 (Gavazzi et al. 2001c; Deharveng et al. 2002), the radio galaxy M87, VCC 1297 (the highest surface brightness galaxy in the sample of Gavazzi et al. (2005)) and VCC 1146. Beside its extremely high surface brightness, making VCC 1297 a non standard object, we do not have any evidence indicating a peculiar star formation history or present nuclear activity in VCC 1297 and VCC 1146 that could justify their exclusion. 76 6. Unveiling the evolution of early type galaxies with GALEX. escent, bulge-dominated Sa spiral disks, that have structural (concentration indices and light profiles) or population properties (colors and spectra) similar to ellipticals and lenticulars (Scodeggio et al. 2002). The monotonic increase of the (N U V − V ) and (N U V − H) colors with luminosity, similar to the one observed in the visible bands by Ferguson (1994) strongly suggests that both in dwarfs and giant systems the NUV 2310 Å flux is dominated by the same stellar population (main sequence low mass stars) emitting at longer wavelengths. On the contrary the different behaviour of the (F U V − V ) and (F U V − H) colors with luminosity, and the clear dichotomy observed in the (F U V − N U V ) vs LH CMR strongly support a different origin for the FUV emission in dwarf and giant systems. The reddening of the UV color index with luminosity observed in dwarf ellipticals, similar to the one observed in late type galaxies, indicates that the UV spectral energy distribution of low mass early type galaxies is shaped by the contribution of young stellar populations. This is shown in Fig.6.3 where the available optical spectra for our sample of dEs are shown. It clearly emerges that UV bluer systems have emission lines or strong Balmer line in absorptions witnessing present or recent star formation activity. Moreover at increasing luminosity their (F U V − N U V ) color index reddens as the optical colors confirming that in these systems the FUV emission is dominated by the contribution of young main sequence stars. This is not the case for giant early type systems: the plateau observed in the FUV-optical CMRs and the blueing of the (F U V − N U V ) color with luminosity (i.e. the UV upturn) suggest that far ultraviolet emission comes from low mass old post asymptotic giant branch stars. This is also confirmed in Figs.6.5 and 6.4 where the optical spectra available for ellipticals and S0s in our sample are presented: as expected, all the spectra are dominated by the contribution of the old stellar populations. Moreover the observed trend between (F U V − N U V ) and the metallicity sensitive Mg2 index, reproduced by models (Bressan et al. 1994; Yi et al. 1998), confirms the early IUE result of Burstein et al. (1988). Conversely Rich et al. (2005) did not find any correlation between the color index (F U V −r) and Mg2 nor with the velocity dispersion σ in a large sample of SDSS early-type galaxies observed by GALEX. Their lack of correlation might derive from insufficient dynamic range in Log σ (2.1-2.4 km s−1 ) and Mg2 (0.18-0.30). The blueing of the UV color index with luminosity, metallicity and velocity dispersion indicates that the UV upturn is more important in massive, metal rich systems. This is consistent with stellar population models which predict that the strength of the UV upturn is mainly driven by stellar metallicity. The accurate morphological classification in our sample allow us to discriminate between E and S0s and to study separately the two populations. The higher dispersion in the (F U V −N U V ) vs. LH relation observed for the lenticulars compared to the extremely tight one for ellipticals (see Table 1), bears witness to a different evolutionary history for the two Hubble types: while cluster ellipticals represent an homogeneous population, S0s are a heterogeneous class probably formed by different independent 6.4. Discussion and conclusion 77 Figure 6.3: The relationship between the UV color index (F U V − N U V ) and the total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for dwarf ellipticals are presented. 78 6. Unveiling the evolution of early type galaxies with GALEX. Figure 6.4: The relationship between the UV color index (F U V − N U V ) and the total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for ellipticals are presented. 6.4. Discussion and conclusion 79 Figure 6.5: The relationship between the UV color index (F U V − N U V ) and the total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for lenticulars are presented. 80 6. Unveiling the evolution of early type galaxies with GALEX. physical mechanisms (see also Chapters 9 and 11), and with various star formation histories as also determined from kinematic and spectroscopic observations (Dressler & Sandage 1983; Neistein et al. 1999; Hinz et al. 2003). Using the available optical spectra we investigate the presence of residual star formation still present in our sample of giant early type galaxies. Only one S0 galaxy, VCC1003, shows a mild residual star formation activity (Hα in emission), while three ellipticals (VCC881,M87,VCC1619) and three S0s (VCC1030,VCC1062,VCC1253) have [NII] in emission and Hα in absorption, a typical feature of low ionization active galactic nuclei. This suggest that the difference observed between ellipticals and S0s cannot be ascribed to recent episodes of star formation but probably resides on their different past star formation history. Combining this result with the one obtained in Chapter 4, we can conclude that, at low UV luminosities, the significant contribution of giant early type systems to the ultraviolet luminosity function must be ascribed not to young stellar populations but to old low mass post-AGB stars. The newest result of this Chapter, shown in Fig. 6.2, addresses the question raised by O’Connell (1999) concerning the dependence of the UV properties on galaxy morphology. We have shown that a dichotomy exists between giant and dwarf ellipticals and, to a lesser extent, between ellipticals and lenticulars. The opposite behavior (reddening of the UV color index with luminosity) of dwarfs with respect to giants, similar to that observed for spirals, indicates that the UV spectra of low luminosity objects are shaped by the contribution of young stars, thus are more sensitive to the galaxy’s star formation history than to the metallicity. This implies that the stellar population of dwarfs has been formed in discrete and relatively recent episodes, as observed in other nearby objects (Grebel 2000). More evidences are building up that mass drives the star formation history in hot systems (Trager et al. 2000; Gavazzi et al. 2002a; Caldwell et al. 2003; Poggianti 2004b) as in rotating ones (Gavazzi et al. 1996; Boselli et al. 2001) and that the stellar population of massive ellipticals is on average older than that of dwarfs. Chapter 7 UV dust attenuation in normal star forming galaxies 7.1 Introduction The use of ultraviolet emission in order to study the properties of star forming galaxies is not an easy a rapid task. The presence of dust in galaxies represents one of the major obstacles complicating a direct quantification of the star formation activity in local and high redshift galaxies. Absorption by dust grains reddens the spectra at short wavelengths completely modifying the spectral energy distribution of galaxies. Since the UV radiation is emitted by young stars (t < 108 yr) that are generally more affected by attenuation from surrounding dust clouds than older stellar populations, rest-frame UV observations can lead to incomplete and/or biased reconstructions of the star formation activity and star formation history of galaxies affected by dust absorption, unless proper corrections are applied. In recent years our understanding of dust attenuation received a tremendous impulse from studies of local starburst galaxies (i.e.Calzetti et al. 1994; Heckman et al. 1998; Meurer et al. 1999; Calzetti 2001; Charlot & Fall 2000), that were based on three indicators: the ratio of the total infrared to far-ultraviolet emission (LT IR /LF U V ), the ultraviolet spectral slope β (determined from a power-law fit of the form f ∼ λβ to the UV continuum spectrum in the range 1300 and 2600 Å, Calzetti et al. 1994) and the Balmer decrement. The total-IR (TIR) to UV luminosity ratio method (i.e. Buat 1992; Xu & Buat 1995; Meurer et al. 1995, 1999) is based on the assumption that a fraction of photons emitted by stars and gas are absorbed by the dust. The dust heats up and subsequently re-emits the energy in the mid- and far-infrared. The amount of UV attenuation can thus be quantified by means of an energy balance. This method is considered the most reliable estimator of the dust attenuation in star-forming galaxies because it is almost completely independent of the assumed 81 82 7. UV dust attenuation in normal star forming galaxies extinction mechanisms (i.e. dust/star geometry, extinction law, see Buat & Xu 1996; Meurer et al. 1999; Gordon et al. 2000; Witt & Gordon 2000). When the spectrum is dominated by a young stellar population the ultraviolet spectral slope β, is found to have a weak dependence on metallicity, IMF, and star formation history (Leitherer & Heckman 1995). Thus the difference between the observed β and the one predicted by models can be entirely ascribed to dust attenuation (Meurer et al. 1999). However in systems with no or mild star formation activity the UV spectral slope can be strongly contaminated by the old stellar populations, whose contribution increases β (flattens the UV continuum, Boissier et al. 2005). Thus the spectral slope of mildly star forming systems could be intrinsically different from the one of starburst galaxies, even in the absence of dust attenuation (Kong et al. 2004). Meurer et al. (1999) have shown that in starburst galaxies the total far-infrared to ultraviolet luminosity ratio correlates with the ultraviolet spectral slope, β (commonly referred to as the IRX-UV relation). They pointed out that this relation allows reliable estimates of the attenuation by dust at ultraviolet wavelengths based on β. The Balmer decrement gives an estimate of the attenuation of ionized gas and not of the stellar continuum as in the previous two methods. It is based on the comparison of the observed Hα/Hβ ratio with its predicted value (2.86 for case B recombination, assuming an electronic density ne ≤ 104 cm−3 and temperature ∼ 104 K; e.g., Osterbrock 1989). Calzetti et al. (1994) found a significant correlation between the ultraviolet spectral slope β and the Balmer decrement Hα/Hβ. Starting from this empirical relation they obtained an attenuation law (known as the Calzetti attenuation law) often adopted to correct UV observations for dust attenuation in absence of both far-infrared observations and estimates of the ultraviolet spectral slope (Steidel et al. 1999; Glazebrook et al. 1999). Unfortunately the above empirical relations have been established only for starburst galaxies and they seem not to hold for normal star forming galaxies. Recently, Bell (2002) suggested that quiescent galaxies deviate from the IRX-UV relation of starburst galaxies, because they tend to have redder ultraviolet spectra at fixed total far-infrared to ultraviolet luminosity ratio. Kong et al. (2004) confirmed this result and interpreted the different behaviour of starbursts and normal galaxies as due to a difference in the star formation histories. They proposed that the offset from the starburst IRX-UV relation can be predicted using the birthrate parameter b (e.g. the ratio of the current to the mean past star formation activity). However an independent observational confirmation of the correlation between the distance from the starburst IRX-UV relation and the birthrate parameter has not been obtained so far (Seibert et al. 2005). Even the Calzetti law does not seem to be universal. Buat et al. (2002) showed that for normal star forming galaxies the attenuation derived from the Calzetti law is ∼0.6 mag larger than the one computed from the FFIR /FUV ratio and their result has been recently confirmed by Laird et al. (2005). Why do normal star-forming galaxies behave differently from starbursts? Do normal 7.1. Introduction 83 galaxies follow different empirical relations that can be exploited to correct for dust attenuation in absence of far infrared observations? If this is the case, is there a transition between starburst and normal galaxies? Which physical parameters drive it? Answering these questions will be important for a better understanding of the interaction of dust and radiation specifically in nearby dusty star forming galaxies, but it also has direct consequences for our understanding and interpretation of galaxy evolution in a general context. Firstly it seems mandatory to characterize the dust attenuation properties of normal galaxies, to compare them with the ones of starbursts and to derive new recipes for the UV dust attenuation correction. This topic came once again to the fore with the launch of the Galaxy Evolution Explorer (GALEX). This satellite is delivering to the community an unprecedented amount of UV data on local and high redshift galaxies that require corrections for dust attenuation but currently lack far-infrared rest-frame data. The time is ripe to explore new methods for correction of these data, that might provide new insights on galaxy evolution. Whenever they can be combined with other data, GALEX observations provide the best available ultraviolet data for studying the dust attenuation properties of galaxies. Multiwavelength photometric and spectroscopic observations are in fact mandatory in order to: determine metallicity, ionized gas attenuation (A(Hα)), luminosity and mass, test the validity of the relations followed by starbursts (Heckman et al. 1998), explore relations that might prove useful to correct ultraviolet magnitudes and to compare them with various models of dust attenuation. Recent extensive spectroscopic and photometric surveys, like the Sloan Digital Sky Survey (SDSS, Abazajian et al. 2005) and the Two Degree Field Galaxy Redshift Survey (2dF, Colless et al. 2001) have opened the path to studies of fundamental physical parameters based on enormous datasets. However, spectroscopic observations of nearby galaxies suffer from strong aperture effects, making these datasets not ideal for the purpose of the present investigation. In fact, Kewley et al. (2005) have recently shown that aperture effects produce both systematic and random errors on the estimate of star-formation, metallicity and attenuation. To reduce at least the systematic effects they suggest selecting only samples with fibres that capture > 20% of the light. This requires z > 0.04 and z > 0.06 for SDSS and 2dF respectively: too distant to detect both giant and dwarf star forming systems with GALEX and IRAS. Although significantly smaller than the SDSS, the dataset we have been building up over the last 10 years with data taken over a large stretch of the electromagnetic spectrum for few thousand galaxies in the local universe (worldwide available from the site GOLDMine; Gavazzi et al. 2003a) turns out to be appropriate for the purposes of the present investigation. It includes drift-scan mode integrated spectra, narrow band Hα and broad band optical and near-infrared imaging for a volume limited sample of nearby galaxies in and outside rich clusters. The combination of GALEX and IRAS observations with these ancillary data allows us to study the dust attenuation properties in a sizable sample of normal star forming galaxies not suffering from the 84 7. UV dust attenuation in normal star forming galaxies aperture bias and to compare observations with model predictions. In this chapter I investigate the relations between dust attenuation and global galaxy properties and compare them with the ones observed in starburst galaxies. The aim of this work is to provide some empirical relations based on observable quantities (thus model independent) suitable for deriving dust attenuation corrections when far infrared data are not available. For this reason all relations obtained throughout this chapter will be given as a function of LT IR /LF U V , the observable that we consider the best dust attenuation indicator. We choose not to transform LT IR /LF U V into a (model dependent) estimate of the far ultraviolet extinction A(F U V ), leaving the reader free to choose his/her preferred dust model (i.e. Meurer et al. 1999; Buat et al. 1999, 2002, 2005; Gordon et al. 2000; Panuzzo et al. 2003; Burgarella et al. 2005, Inoue et al. in preparation). We assume that quantities are related linearly and residual plots are presented in order to test the validity of this hypothesis. Moreover, since we are looking for new recipies to estimate the LT IR /LF U V ratio, this quantity has to be considered as the dependent variable, implying the use of an unweighted simple linear fit to estimate the best fitting parameters (Isobe et al. 1990). 7.2 7.2.1 The Data The optically-selected sample The analysis presented in this work is based on an optically selected sample of latetype galaxies (later than S0a) including giant and dwarf systems extracted from the Virgo Cluster Catalogue (VCC, Binggeli et al. 1985) and from the CGCG catalogue (Zwicky et al. 1961). The data include ∼ 300 square degrees covering most of the Virgo, Abell1367 and Abell262 clusters, the southwest part of the Coma cluster and part of the Coma-A1367 supercluster (11h30m < R.A. < 13h30m; 18◦ < decl. < 32◦ ) observed in spring 2004 as part of the All-sky Imaging Survey (AIS) and of the Nearby Galaxy Survey (NGS) carried out by GALEX in two UV bands: FUV (λeff = 1530Å, ∆λ = 400Å) and NUV (λeff = 2310Å, ∆λ = 1000Å). Details of the GALEX instrument and characteristics can be found in Martin et al. (2005) and Morrissey et al. (2005). Our sample has the quality of being selected with the criterion of optical completeness. All galaxies brighter than a threshold magnitude are selected in all areas. In Coma-A1367 supercluster and A262 cluster all galaxies brighter than mp =15.7 were selected from the CGCG catalogue (Zwicky et al. 1961). The Virgo region contains all galaxies brighter than mp =18 from the VCC catalogue (Binggeli et al. 1985). We thus consider our sample an optically selected, volume limited sample. We include in our analysis all late-type galaxies, detected in both NUV and FUV GALEX bands and in both 60 µm and 100 µm IRAS bands (157 objects). When- 7.2. The Data 85 ever available, we extracted UV fluxes from the deep NGS images, obtained with a mean integration time of ∼ 1500 sec, complete to mAB ∼ 21.5 in the NUV and FUV. Elsewhere UV fluxes have been extracted from the shallower AIS images (∼ 70 sq. degrees), obtained with a mean integration time of ∼ 100 sec, complete to mAB ∼ 20 in both the FUV and NUV bands. All UV images come from the Internal Data Release v1 (IR1.0). UV fluxes were obtained by integrating GALEX images within elliptical annuli of increasing diameter up to the optical B band 25 mag arcsec−2 isophotal radii, consistently with the optical and near-IR images. Independent measurements of the same galaxies obtained in different exposures give consistent photometric results within 10 % in the NUV and 15% in the FUV in the AIS, and a factor of ∼ two better for bright (NUV ≤16) galaxies. The uncertainty in the UV photometry is on average a factor of ∼ 2 better in the NGS than in the AIS, particularly for faint objects. The typical uncertainty in the IRAS data is 15% (Boselli et al. 2003a). UV and far-infrared data have been combined to multifrequency data. These are optical and near-IR H imaging (mostly from Gavazzi et al. 2000, 2005; Boselli et al. 2003a), optical drift-scan spectra (Gavazzi et al. 2004; Gavazzi et al. in prep.) and Hα imaging (Boselli & Gavazzi 2002; Boselli et al. 2002a; Gavazzi et al. 1998, 2002b; IglesiasPáramo et al. 2002; Gavazzi et al. in prep.), great part of which are available from the GOLDMine galaxy database (Gavazzi et al. 2003a) (http://goldmine.mib.infn.it). From the 157 galaxies selected we exclude Active Galactic Nuclei (AGN). AGNs have been selected using either the classification provided by NED, if available, or by inspection to the integrated spectra of Gavazzi et al. (2004): we exclude galaxies with log([OIII]/Hβ) > 0.61/(log([N II]/Hα) − 0.05) + 1.3 (Kauffmann et al. 2003a). This criterion reduces the sample to 128 galaxies, spanning a range of six magnitudes in B band (-22< MB <-16) and of three orders of magnitude in mass1 (9 < M < 12 M ). Unfortunately ancillary data are not available for all galaxies observed by GALEX, we thus further divided the data in two subsamples. Sixty six galaxies in the primary sample have all the necessary complementary data (e.g. Hα photometry, Hα/Hβ ratio, metallicity, H-band photometry; see Gavazzi et al. 2000, 2002a,b, 2004 for the selection criteria adopted in each survey). The remaining 62 galaxies form the secondary sample. We cannot exclude a possible contamination of AGN in the secondary sample, since no spectra are available for these objects. In all figures objects belonging to the primary sample will be indicated with filled circles while the secondary sample as empty circles. Since only galaxies belonging to the primary sample are present in all the plots analyzed in the presented work, all correlations will be quantified using only the primary sample. Data from UV to near-IR have been corrected for Galactic extinction according to Burstein & Heiles (1982). We assume a distance of 17 Mpc for the members of Virgo Cluster A, 22 Mpc for Virgo Cluster B, and 32 Mpc for objects in the M and W clouds (Gavazzi et al. 1999a). 1 Computed using the relation between LH and M by Gavazzi et al. (1996) 86 7. UV dust attenuation in normal star forming galaxies Members of the Cancer, A1367, and Coma clusters are assumed to lie at distances of 65.2, 91.3, and 96 Mpc, respectively. Isolated galaxies in the Coma supercluster are assumed at their redshift distance, adopting H0 = 75 km s−1 Mpc−1 . 7.2.2 The starburst sample In order to compare the properties of our sample with starbursts, we compile a dataset of starburst galaxies observed by IUE from the sample of Calzetti et al. (1994). We consider 29 galaxies, excluding AGNs and galaxies that have not been observed by IRAS at 60 or 100 µm. Complementary data such as FIR, Hα fluxes and Balmer decrements are taken from Calzetti et al. (1995), metallicities come from Heckman et al. (1998) and H-band photometry (available only for 18 galaxies) from (Calzetti 1997). Excluding the far infrared fluxes, all these quantities are obtained within an apertures of ∼ 20 × 10arcsec2 , consistent with IUE observations Calzetti et al. (1994). Thus we stress that aperture effects could strongly affect any comparison with normal galaxies for which all data are homogeneously integrated values. First of all, if the UV emission is more extended than IUE field of view the LT IR /LF U V ratio is overestimated2 . In addition, even when physical quantities are obtained in the same IUE apertures, the presence of age and metallicity gradients in galaxies makes not trivial any comparison with the integrated values obtained for normal star forming galaxies (Kewley et al. 2005). All the observables, but the ultraviolet spectra slope β, are calibrated in a consistent way with our sample of normal galaxy. The ultraviolet spectral slope of starbursts is obtained by fitting IUE spectra (Calzetti et al. 1994), while for GALEX observations it comes from the FUV-NUV color index (see next Section). However, as shown by Kong et al. (2004), these two calibrations are consistent each other and do not introduce any systematic difference between the two samples. 7.3 The LT IR/LF U V − β relation for normal starforming galaxies Meurer et al. (1999) have shown that the ratio of far infrared to far ultraviolet luminosity tightly correlates with the UV colors of starburst galaxies. This relation, known as the infrared excess-ultraviolet (IRX-UV) relation, is often presented as β vs. LT IR /LF U V relation. As discussed in the introduction, we will refer to the LT IR /LF U V ratio as the best indicator of UV dust attenuation and we will calibrate on it all the following relations. In order to determine the dust emission, we compute the total 2 However Meurer et al. (1999) argued that the majority of UV flux for their starburst sample lies within the IUE aperture 7.3. The LT IR /LF U V − β relation for normal star-forming galaxies 87 infrared flux emitted in the range 1-1000 µm, following Dale et al. (2001): f60 )+ f100 f60 2 f60 3 +0.7281 × log( ) + 0.6208 × log( ) + f100 f100 f60 4 ) +0.9118 × log( f100 log(fT IR ) = log(fF IR ) + 0.2738 − 0.0282 × log( (7.1) where fF IR is the far-infrared flux, defined as the flux between 42 and 122 µm (Helou et al. 1988): fF IR = 1.26 × (2.58 × f60 + f100 ) × 10−14 [Wm−2 ] (7.2) and f60 and f100 are the IRAS fluxes measured at 60 and 100 µm (in Jansky). The total infrared luminosity is thus: LT IR = 4πD 2 fT IR (7.3) The β parameter as determined from GALEX colors is very sensitive to the galaxy star formation history (see for example Calzetti et al. 2005). For this reason we assume β as defined by Kong et al. (2004): log(fF U V ) − log(fN U V ) = −0.182 = 2.201 × (F U V − N U V ) − 1.804 β= (7.4) where fF U V and fN U V are the near and far ultraviolet observed fluxes respectively (in erg cm2 s−1 Å−1 ), and FUV and NUV are the observed magnitudes. The relationship between the ratio of total infrared luminosity (LT IR ) obtained from (7.1) to the far-ultraviolet fluxes and the UV spectral slope β (or the FUV-NUV color) for our sample of nearby star forming galaxies is given in Fig.7.1. Several functional forms of the LT IR /LF U V − β relation can be found in the literature (i.e. Meurer et al. 1999; Kong et al. 2004); we simply adopt a linear fit: log(LT IR /LF U V ) = a × β + b. This functional form is consistent with other previously proposed for β > −2, while it diverges for β < −2. Since the majority of normal and starbursts galaxies have β > −2 our choice is justified. This represents the simplest and less parameter dependent way to study the relation between two quantities.3 We find a strong correlation (Spearman correlation coefficient rs ∼0.76 for the primary sample and rs ∼0.65 for the secondary sample, both corresponding to a probability P (rs ) >99.9% 3 We tested this hypothesis fitting our data with functional forms similar to the ones proposed by Meurer et al. (1999) and Kong et al. (2004): no significative improvement in the scatter of this relation is obtained. 88 7. UV dust attenuation in normal star forming galaxies Figure 7.1: Ratio of the total infrared to far ultraviolet luminosity as a function of the ultraviolet spectral slope (lower x-axis) and the FUV-NUV color (upper x-axis). Open circles indicates our secondary sample while filled circles represent the primary sample. The dashed line represents the best linear fit to starburst IRX-UV relation. The solid line indicates the best bisector linear fit for our primary sample. The stars indicate the sample of IUE starbursts. Mean error bars for the plotted data are shown in the lower right corner, in this and subsequent figures. The residuals from the best linear fit for normal galaxies are shown in the bottom panel. 7.3. The LT IR /LF U V − β relation for normal star-forming galaxies 89 that the two variables are correlated) between the total infrared to far ultraviolet ratio and the spectral slope, but significantly different from the one observed for starburst galaxies (dashed line in Fig.7.1; Meurer et al. 1999). A χ2 test rejects at a confidence level higher than 99.9%, the hypothesis that the two samples follow the same relation. The best linear fit for our primary sample (solid line in Fig.7.1) is: log( LT IR ) = (0.70 ± 0.06) × β + (1.30 ± 0.06) LF U V (7.5) The uncertainty in the estimate of the LT IR /LF U V using equation (7.5) is ∼ 0.26±0.02 dex for the primary sample but it increases to ∼ 0.35 ± 0.03 dex, if we consider the whole sample (e.g. primary and secondary samples), consistent with the mean uncertainty observed for starburst galaxies (Meurer et al. 1999). A large contribution (∼ 0.21±0.02 dex) to the observed scatter in Eq.(7.5) is due to the uncertainty on the estimate of LT IR /LF U V and β. This result confirms once more that the LT IR /LF U V −β relation for normal galaxies deviates from the one observed for starbursts, as pointed out by previous studies of nearby galaxies (i.e. Bell 2002; Kong et al. 2004; Boissier et al. 2005; Buat et al. 2005; Seibert et al. 2005; Burgarella et al. 2005, Boissier et al. in prep.) and individual HII regions in nearby galaxies (Calzetti et al. 2005). 7.3.1 The dependence on the birthrate parameter What physical mechanisms drive the difference observed in the LT IR /LF U V − β between normal star forming galaxies and starbursts? Recently Kong et al. (2004) interpreted the offset as an effect of the different star formation history experienced by galaxies and proposed that the distance from the starburst IRX-UV can be predicted using the birthrate parameter b (e.g. the ratio of the current to the mean past star formation activity, Kennicutt et al. 1994). In order to test if the perpendicular distance dS from the LT IR /LF U V − β relation for starbursts correlates with the star formation history of normal galaxies, we compute the birthrate parameter following Boselli et al. (2001): SF Rt0 (1 − R) b= (7.6) LH (Mtot /LH )(1 − DMcont ) where R is the fraction of gas that stellar winds re-injected into the interstellar medium during their lifetime (∼ 0.3, Kennicutt et al. 1994), t0 is the age of the galaxy (that we assume ∼12 Gyr), DMcont is the dark matter contribution to the Mtot /LH ratio at the optical radius (assumed to be 0.5; Boselli et al. 2001). We compute the H-band luminosity following Gavazzi et al. (2002a): log LH = 11.36 − 0.4 × H + 2 × log(D) [L ] 90 7. UV dust attenuation in normal star forming galaxies Figure 7.2: Relation between the birthrate parameter computed from the Hα emission, and the distance from the LT IR /LF U V − β relation for starbursts. The solid line represents the best linear fit. where D is the distance to the source (in Mpc), and the SFR from the Hα luminosity (corrected for [NII] contamination and for dust extinction using the Balmer decrement) following Boselli et al. (2001): SF R = LHα 1.6 × 1041 [M /yr] (7.7) Fig.7.2 shows the relation between the birthrate parameter (eq.7.6) and the distance from the LT IR /LF U V − β relation for starburst galaxies. The two quantities are correlated (rs ∼0.40, corresponding to a correlation probability P (rs ) ∼99.8%) but with a large scatter. Given the value of observational uncertainties, it is not worth trying to use the observed trend to reduce the dispersion in the LT IR /LF U V − β relation for normal galaxies. This result confirms that part of the dispersion in the LT IR /LF U V − β relation for normal star forming galaxies appears an effect of the different star formation history experienced by galaxies, as proposed by Kong et al. (2004). 7.4 7.4.1 A(Hα) Estimate of A(Hα) The attenuation in the Balmer lines can be deduced from the comparison of the observed ratio LHα /LHβ with the theoretical value of 2.86 obtained for the recombination case B, an electronic density ne ≤ 104 cm−3 and temperature ∼ 104 K. The 7.4. A(Hα) 91 variation of this value with density its negligible and with temperature is ≤5% (in the range between 5000 K and 20000 K, Caplan & Deharveng 1986). The underlying absorption was deblended from the Hβ emission line using a multiple component fitting procedure. To do this the emission line is measured and subtracted from the spectra. The resulting absorption line is also measured with respect to a reference continuum. These two measurements are used as first guess in a fitting algorithm which fits jointly the emission and absorption lines to the reference continuum. For objects whose Hβ is detected in emission but the deblending procedure is not applied (no absorption feature is evident) a mean additive correction for underlying absorption equal to -1.8 in flux and -1.4 Åin EW is used. These values correspond to the fraction of the (broader) absorption feature that lies under the emission line. We adopt a dust screen geometry and the Milky Way extinction curve (e.g. Kennicutt 1983; Calzetti et al. 1994). Whereas varying the extinction curves has negligible effects in the visible, the dust screen assumption seems to under-estimate the extinction by ∼0.2 mag compared with the amount deduced from the measurements of the thermal radio continuum (Caplan & Deharveng 1986; Bell & Kennicutt 2001). We do not apply any correction for Hα underlying absorption (Charlot & Longhetti 2001). However, since all the objects have EW (Hα + [N II]) > 3Å, the underestimate in the value of A(Hα) is negligible. In fact no change (at a 99% significance level) is observed comparing the best fits obtained in this work and the ones obtained adding to the Hα the same fixed underlying absorption used for Hβ when the underlying is not detected. We assume that the errors on A(Hα) are mainly due to the uncertainty on the Hβ flux. These errors represent in fact the lower limits because we do not account for the uncertainty introduced by the fitting of the lines. They range from 0.01 to 0.43 mag and are found strongly anti-correlated with EW(Hβ) (see Gavazzi et al. 2004). Adopting the definition of the Balmer decrement as in Gavazzi et al. (2004): C1(Hβ) = 1 log( 2.86 × LHα ) LHβ 0.33 (7.8) Since the A(Hα) attenuation is: 1 LHα 1 ln( × ) eβα − 1 2.86 LHβ (7.9) 1 × 0.33 × C1(Hβ) ln(10) eβα − 1 (7.10) A(Hα) = 1.086 From (7.8) and (7.9) we obtain: A(Hα) = 1.086 92 7. UV dust attenuation in normal star forming galaxies and assuming a galactic extinction law (eβα = 1.47) we derive: A(Hα) = 1.756 × C1(Hβ) (7.11) A(Hα) = 0.85 mag is obtained on average, consistent with previous studies (e.g. Kennicutt 1983, 1992; Thuan & Sauvage 1992; Kewley et al. 2002). Eleven galaxies have Hβ undetected in emission but the underlying stellar absorption is clearly detected. For them we derive a 3×σlower limit to the Hβ flux (fHβ ) using (Gavazzi et al. 2004): fHβ < 3 × rms(4500−4800) × Hα(HW HM ) (7.12) assuming that Hα and Hβ emission lines have similar HWHM (Half Width Half Maximum). As shown in Eq.(7.8) a change in the theoretical value of the LHα /LHβ ratio would only produce a small (≤5%) constant over (or under) estimate of the ionized gas attenuation, thus leaving unchanged the shape and dispersions of the observed relations, only affecting the values of the best fitting parameters. 7.4.2 The β-A(Hα) relation Calzetti et al. (1994) found a strong relationship between the ultraviolet spectral slope β and the Balmer decrement Hα/Hβ. For our starburst sample these two quantities are correlated (rs ∼0.81) as follows (see also blue stars in Fig.7.3): β = (0.75 ± 0.10) × A(Hα) − (1.80 ± 0.13) (7.13) This empirical relation was used by Calzetti et al. (1994) to deduce an attenuation law (the Calzetti law), often applied to high redshift galaxies (i.e. Steidel et al. 1999; Glazebrook et al. 1999). Contrary to the LT IR /LF U V − β relation the Calzetti law has not yet been tested for a sample of normal star forming galaxies. Buat et al. (2002) showed that for normal star forming galaxies the attenuation derived from the Calzetti law is ∼0.6 larger than the one computed from F IR/U V ratio. This result has been recently confirmed by Laird et al. (2005) on star forming galaxies at z ∼1. In order to check the Calzetti law on our sample we use the measure of the Hα/Hβ described in the previous subsection. Fig. 7.3 shows the relation between β and A(Hα) for our sample (empty and filled circles). For the primary sample we obtain rs ∼0.58 (P (rs ) >99.9%) and: β = (0.37 ± 0.07) × A(Hα) − (1.15 ± 0.08) (7.14) flatter than for starburst galaxies (see Fig.7.3). At low A(Hα) normal galaxies show on average a less steep ultraviolet spectral slope than starbursts. In addition normal galaxies with the same value of β span a range of ∼ 1 mag in A(Hα). At higher 7.4. A(Hα) 93 Figure 7.3: The relation between the ultraviolet spectral slope β and the Hα attenuation obtained from the Balmer decrement. Symbols are as in Fig.7.1. Solid line represents the best linear fit to our primary sample (equation 7.14) while the dashed line indicate the best-fit for starburst galaxies obtained by Calzetti et al. (1994) (equation 7.13). Arrows indicate galaxies for which the value of A(Hα) is a lower limit of the real value (i.e. Hβ undetected). The residuals from the best linear fit for normal galaxies are shown in the bottom panel. 94 7. UV dust attenuation in normal star forming galaxies attenuation the two samples appear consistent. Our result suggest that the Calzetti law cannot be applied to normal galaxies. On the contrary, the relation between β and A(Hα) for normal galaxies, could be used to obtain a new attenuation law. 7.5 7.5.1 Relations between dust attenuation and global properties. Metallicity Heckman et al. (1998) have shown that the ultraviolet spectral slope and metallicity of starbursts are well correlated. To determine the metal content of our galaxies we average five different empirical determinations based on the following line ratios: R23 ≡ ([OII]λ3727 + [OIII]λ4959, 5007)/Hβ (Zaritsky et al. 1994; McGaugh 1991), [NII]λ6583/[OII]λ3727 (Kewley & Dopita 2002), [NII]λ6583/Hα (van Zee et al. 1998) and [OIII]λ5007/[NII]λ6583 (Dutil & Roy 1999). The mean uncertainty in the abundances is 0.10dex. In Fig. 7.4 we study the relationship between the gas metallicities and the LT IR /LF U V ratio (left-panel) and β (right-panel) for normal star forming and starburst galaxies. For normal galaxies the LT IR /LF U V ratio correlates (rs ∼0.59, P (rs ) >99.9%) with the gas abundance: log( LT IR ) = (1.37 ± 0.24) × 12 + log(O/H) − (11.36 ± 2.11) LF U V (7.15) with a dispersion of ∼ 0.35 ± 0.03 in log(LT IR /LF U V ). As for the LT IR /LF U V − β relation normal galaxies differ from starbursts. At comparable metallicity normal galaxies show a lower LT IR /LF U V (lower attenuation) than starbursts, in agreement with the recent result by Boissier et al. (2004) who studied radial extinction profiles of nearby late-type galaxies using FOCA and IRAS observations. Unexpectedly we find however that normal star forming galaxies follow exactly the same (significant, rs ∼ 0.58, P (rs ) >99.9%) relationship between metallicity and ultraviolet spectral slope β determined for starbursts by Heckman et al. (1998) (see right panel of Fig.7.4). This might indicate that even though a normal and a starburst galaxy with similar gas metallicity have similar UV spectral slopes, they suffer from a significantly different dust attenuation, perhaps suggesting a different dust geometry (Witt & Gordon 2000). However we stress that this effect might occur due to aperture effects in the IUE data: while β is not significantly contaminated by aperture effects, the LT IR /LF U V ratio could be overestimated producing the observed trend (the total infrared luminosity is obtained by integrating the IRAS counts over the full galaxy extension, while the ultraviolet one is taken from IUE’s significantly smaller aperture 20 × 10 arcsec2 ). This idea could be supported by the correlation (rs ∼ 0.49, 7.5. Relations between dust attenuation and global properties. 95 Figure 7.4: Relation between gas metallicity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig.7.1. The solid lines show the best linear fit for our primary sample. The residuals from the best linear fits for normal galaxies are shown in the upper panels. Figure 7.5: Relation between the galaxy size and the LT IR /LF U V ratio for starburst (left panel) and normal galaxies (right panel). Symbols are as in Fig. 7.1. Mean values and uncertainties in bins of 0.30 log(Diameter) are given. 96 7. UV dust attenuation in normal star forming galaxies P (rs ) >99.9% see Fig.7.5) observed between the starbursts’ optical diameters and the LT IR /LF U V ratio, completely absent in our sample of normal galaxies (rs ∼ 0.006, P (rs ) ∼25%). GALEX observations of starburst galaxies will rapidly solve this riddle. Dust to Gas ratio The correlation between attenuation and metallicity can be interpreted assuming that the ultraviolet radiation produced by star forming regions suffers a dust attenuation increasing with the dust to gas ratio, which correlates with metallicity. (e.g. Issa et al. 1990; Inoue 2003). In order to check this hypothesis we compute the dust to gas ratio following Boselli et al. (2002b). In normal galaxies the dust mass is dominated by the cold dust emitting above ∼200 µm. The total dust mass can be estimated provided that the 100-1000 µm far-IR flux and the cold dust temperature are known. Fitting the SEDs of normal galaxies with a modified Planck law ν β Bν (Td ), with β = 2 (Alton et al. 2000), the total dust mass can be determined from the relation (Devereux & Young 1990): Mdust = CSλ D 2 (ea/Tdust − 1) M (7.16) where C depends on the grain opacity, Sλ is the far-IR flux at a given wavelength (in Jy), D is the distance of the galaxy (in Mpc), Tdust is the dust temperature, and a depends on λ. Only IRAS data at 60 and 100 µm are available for our sample and, given the strong contamination of the emission at 60 µm by very small grains, the 60 to 100 µm ratio does not provide a reliable measure of Tdust (Contursi et al. 2001). Tdust determined by Alton et al. (1998) consistently with Contursi et al. (2001), seems to be independent of the UV radiation field, of the metallicity or of the total luminosity (Boselli et al. 2002b). Therefore we will adopt the average value Tdust = 20.8 ± 3.2 K for all our galaxies introducing an uncertainty of ∼50% on the estimate of Mdust (equation (7.16)). We then estimate the dust mass of the sample galaxies using (7.16) with C = 1.27 M Jy−1 Mpc−2 , consistent with Contursi et al. (2001), and a=144 K for Sλ = S100 µm (Devereux & Young 1990). The determination of the dust to gas ratio, in a way consistent with that obtained in the solar neighbourhood, requires the estimate of the gas and dust surface densities, thus of the spatial distribution of dust and gas over the discs. Unfortunately only integrated HI and H2 masses are available for our spatially unresolved galaxies. It is however reasonable to assume that the cold dust and the molecular hydrogen are as extended as the optical disc (Alton et al. 1998; Boselli et al. 2002b). To determine the mean HI surface density we adopt (Boselli et al. (2002b)): log ΣHI = 20.92(±0.17) − 0.65(±0.11) × (def (HI)) cm−2 where def(HI) is the galaxy HI deficiency. Thus the dust to gas ratio is obtained from 7.5. Relations between dust attenuation and global properties. 97 Figure 7.6: Relation between the gas to dust ratio and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. the ratio of the dust surface density to the sum of molecular and neutral hydrogen surface densities. In Fig. 7.6 we compare the relation between the LT IR /LF U V ratio (left panel) and β (right panel) with the dust to gas ratio. The gas to dust ratio barely correlates with the LT IR /LF U V ratio (R∼0.38). Contrary to metallicity, we do not find a significant correlation (R∼0.11) with the ultraviolet spectral slope. This is probably due to the high uncertainty in our estimate of Mdust consequent to assuming the same temperature for all our galaxies (Mdust ∝ ea/Tdust , thus small errors (∼15%) on Tdust propagate onto ∼50% errors on Mdust ). 7.5.2 Luminosity Since it is well known that the metallicity of normal galaxies strongly correlates with galaxy luminosity (e.g. Skillman et al. 1989; Zaritsky et al. 1994) and mass (e.g. Tremonti et al. 2004), it is worth considering the correlation between attenuation and galaxy luminosity. Fig.7.7 shows the relationships between the dust attenuation indicators LT IR /LF U V and β and the H-band luminosity. The infrared to far ultraviolet ratio correlates (rs ∼0.49, P (rs ) >99.9%) with the total H-band luminosity: log( LT IR LH ) = (0.34 ± 0.10) × log( ) − (2.66 ± 0.88) LF U V L (7.17) The dispersion of this relation is ∼ 0.39 ± 0.03 in log(LT IR /LF U V ). Since the H-band luminosity is proportional to the dynamical mass (Gavazzi et al. 1996), this implies a relationship between dust attenuation and dynamical mass. Also in starbursts the 98 7. UV dust attenuation in normal star forming galaxies Figure 7.7: Relation between the H-band luminosity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panel. Figure 7.8: Relation between the TIR+FUV luminosity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. 7.5. Relations between dust attenuation and global properties. 99 total H-band luminosity is correlated (rs ∼0.37, P (rs ) ∼99.5%) with the LT IR /LF U V ratio and the great part of starbursts appear offset (to 99% confidence level) from the relation of normal galaxies. On the contrary, no difference is observed between the two samples in the β-LH plot, in agreement with what observed for metallicity. Finally Fig.7.8 shows the relation between the bolometric luminosity (LTIR + LFUV ) and the dust attenuation, computed assuming that the UV emission is absorbed by dust and emitted in the far infrared. The correlation coefficient (rs ∼0.31, P (rs ) ∼98%) indicates that the two quantities correlate, as for starburst galaxies (Heckman et al. 1998). This is not the case if we examine the relation between the ultraviolet spectral slope β and the bolometric luminosity (Fig.7.8 right panel): while there is no correlation (rs ∼ 0.002, P (rs ) ∼20%) for our sample of normal galaxies, a clear relation (rs ∼ 0.68, P (rs ) >99.9%) holds for starbursts. Starbursts with higher bolometric luminosity (high TIR emission) show lower ultraviolet slope, consistent with the idea that high TIR emission corresponds to high attenuation (low β). 7.5.3 Surface brightness Wang & Heckman (1996) interpreted the increase of dust attenuation with rotational velocity (or mass) as due to the variations in both the metallicity and surface density of galactic disk with galactic size. Fig.7.9 shows the variation of the effective H-band surface brightness (defined as the mean surface brightness within the radius that contains half of the total galaxy light) and the dust attenuation. The two quantities are strongly anti-correlated (rs ∼-0.63, P (rs ) >99.9%): log( LT IR ) = (−0.28 ± 0.04) × µe (H) + (5.92 ± 0.81) LF U V (7.18) with a scatter of ∼ 0.34 ± 0.03 in log(LT IR /LF U V ): ∼1.2σ lower than the value obtained for H-band luminosity and consistent with the one obtained for the gas metallicity. Unfortunately in this case we cannot compare the behaviour of normal galaxies with the one of starbursts due to the lack of an estimate of µe for the starbursts. Does this relation indicate that UV dust extinction depends on the thickness of stellar disk, or does it follows from the correlation between attenuation and star formation surface density? To attack this question we determine the SFR density (defined as the ratio between the SFR determined from Hα (eq.7.7) and optical galaxy area). Fig.7.10 shows the relation between the SFR density and log(LT IR /LF U V ). The two quantities are correlated (rs ∼0.44, P (rs ) >99.9%) with a dispersion of ∼ 0.39 ± 0.03 in log(LT IR /LF U V ), ∼1.2σ larger than the one observed for the mean H-band surface brightness4 . Since the contribution of observational uncertainties to 4 The difference between the two relations does not change if instead of the half-light radius, we use the total radius to estimate µe (H) 100 7. UV dust attenuation in normal star forming galaxies Figure 7.9: Relation between the mean H-band surface brightness (µe ) and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panel. Figure 7.10: Relation between the star formation rate density and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample.The residuals from the best linear fit for normal galaxies are shown in the upper panel. 7.5. Relations between dust attenuation and global properties. 101 the scatter in the two relations is ∼ the same (0.18 ± 0.02), our result might suggest that the UV attenuation is primarily correlated with the thickness of stellar disk, supporting the hypothesis of Wang & Heckman (1996) that both gas metallicity and star surface density are directly connected with the physical properties of dust (i.e. quantity and spatial distribution). 7.5.4 LHα /LF U V ratio Buat et al. (2002) suggested that the LHα /LF U V ratio could be another potential attenuation indicator but they found a scattered correlation between LHα /LF U V and A(F U V ), confirmed by Bell (2002). This correlation is expected since both Hα and UV emission are star formation indicators. The Hα luminosity comes from stars more massive than 10 M and it traces the SFR in the last ≤ 107 yr while the UV luminosity comes from stars of lower mass (M≥ 1.5 M ) and it can be used as an indicator of the SFR in the last ≈ 108 yr. This means that under the condition that the star formation is approximately constant in the last ≈ 108 yr the ratio LHα /LF U V (corrected for attenuation) should be fixed. Thus the ratio between the extinction corrected LHα and the observed LF U V should be a potential attenuation indicator. In Fig.7.11 we analyze the relationship between the dust attenuation and the LHα /LF U V ratio, where LHα is the Hα luminosity corrected for dust attenuation using the Balmer decrement and for the contamination of [NII]. The two quantities turn out to be strongly correlated (rs ∼0.76, P (rs ) >99.9%): log( LHα LT IR ) = (0.84 ± 0.07) × log( ) − (0.59 ± 0.12) LF U V LF U V (7.19) The dispersion around this relation is ∼ 0.24 ± 0.02 in log(LT IR /LF U V ), consistent with the one observed for the log(LT IR /LF U V ) − β relation. The high correlation and low scatter between the two quantities is expected since the two variables are mutually related: the FUV luminosity appears in both axes and LT IR and LHα are known to be correlated (Kewley et al. 2002), explaining why in the left panel of Fig.7.11 starbursts and normal galaxies show the same trend. The right-panel of Fig.7.11 shows the relation between the ultraviolet slope and the LHα /LF U V ratio. In this case starbursts and normal galaxies behave differently: at any given β starbursts have an higher LHα /LF U V than normal galaxies, consistent with what expected for galaxies experiencing a burst of star formation (Iglesias-Páramo et al. 2004). A secure determination of the Balmer decrement for large samples is still a hard task, especially at high redshift, thus we look for a relation similar to Eq.(7.19) using obs the observed Hα luminosity (Lobs Hα ). The LHα /LF U V and log(LT IR /LF U V ) ratios are yet correlated (see Fig.7.12) but the correlation coefficient is lower than the previous 102 7. UV dust attenuation in normal star forming galaxies Figure 7.11: Relation between the Hα and far ultraviolet luminosity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. Hα luminosity is corrected for dust attenuation using the Balmer decrement, while the FUV flux is uncorrected. The solid lines show the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panels. Figure 7.12: Relation between the observed Hα and far ultraviolet luminosity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. Hα luminosity is the observed value not corrected for dust attenuation. The solid lines show the best linear fit for our primary sample.The residuals from the best linear fit for normal galaxies are shown in the upper panels. 7.6. A cookbook for determining LT IR /LF U V ratio 103 case (rs ∼0.49, P (rs ) >99.9%). The best linear fit gives: log( LT IR Lobs ) = (1.10 ± 0.17) × log( Hα ) − (0.59 ± 0.21) LF U V LF U V (7.20) with a mean absolute deviation of ∼ 0.34 ± 0.03 (∼3.3σ higher than for Eq.7.20). 7.6 A cookbook for determining LT IR /LF U V ratio In this chapter we investigated the relations between dust attenuation, traced by the LT IR /LF U V ratio, and other global properties of normal star forming galaxies. Furthermore we compared the dust attenuation in normal and starbursts galaxies using multiwavelength datasets. The amount of dust attenuation is found to correlate with the UV colors, gas metallicity, mass and mean surface density but, generally speaking, differently for normal and starburst galaxies. Determine whether this difference is real or is due to aperture effects requires the analysis of GALEX observations for a sample of starburst galaxies. The dispersion in the LT IR /LF U V − β relation correlates with the birthrate parameter b, suggesting that the observed scatter is, at least partly, due to differences in the star formation history. These results stress that estimating the UV dust attenuation, and consequently the star formation rate of normal galaxies (at high redshift in particular) is highly uncertain (≥50%) when rest-frame far infrared observations are not available. Moreover the sample selection criteria could strongly affect its properties, as recently pointed out by Buat et al. (2005) and Burgarella et al. (2005). They studied the dust attenuation properties and star formation activity in a UV and in a FIR selected sample, showing that the former shows correlations with global galaxy properties, such as mass and bolometric luminosity, that the FIR selected sample does not. Their results stress that the dust attenuation properties are very heterogeneous and that LT IR /LF U V cannot be derived in a robust manner when far infrared observations are not available. However the present investigation has shown that among optically-selected samples of normal galaxies with no nuclear activity a number of empirical relations exists, allowing to derive the LT IR /LF U V ratio (and its uncertainty). Once the attenuation at UV is determined it can be transformed to any other λ, only knowing the shape of the attenuation law and dust geometry (i.e. Calzetti et al. 1994; Gavazzi et al. 2002a; Boselli et al. 2003a). In Table 7.1 we list all the relations, their associated r.m.s., mean absolute deviation from the best fit (m.a.d.)5 and the Spearman correlation coefficient. 5 The mean absolute deviation is less sensitive to the contribution of outliers than p the standard deviation. For a Gaussian distribution the mean absolute deviation (m.a.d.) is ∼ 2/π × (r.m.s.), while it is lower (higher) for a heavier (lighter) tailed distribution. As shown in Table 1 the values 104 7. UV dust attenuation in normal star forming galaxies Before we proceed describing our recipes, we have to investigate whether the scatter in these relations is physical or is only driven by observational uncertainties. In the latter case, in fact, our cookbook would not be very useful, since it would be valid only for observations with the same uncertainties as our datasets. For H-band luminosity, H-band surface brightness, Lobs Hα /LF U V ratio and metallicity the contribution of observational uncertainties to the observed scatter varies from ∼ 18% (r.m.s.∼ 0.17 ± 0.02) for LH to ∼40% (r.m.s.∼ 0.21 ± 0.02) for 12 + log(O/H) and Lobs Hα /LF U V : even accounting for the contribution of measurements errors, the relative difference in the scatter of these relations does not change. On the contrary this confirms that the relation involving LH is the one with the highest ”physical” dispersion, while for the other three relations the scatter is similar. The situation is worse for the relations involving β and the LHα /LF U V ratio: the contribution of observational errors is ∼70-76% (∼ 0.21 ± 0.02). Thus it is impossible to determine which of these two relations has the lowest scatter and represents the best way to estimate dust attenuation without far infrared observations. We can conclude that observational errors could account for the difference scatter observed in the relations involving β and the LHα /LF U V ratio, but not for the difference observed in all the other relations. Our results can thus be used to suggest different ways to correct for UV dust attenuation. Ia) The LT IR /LF U V − β relation still represents one of the best way to quantify dust attenuation. The uncertainty in the value of log(LT IR /LF U V ) is ∼ 0.26 ± 0.03. Ib) If the UV spectral slope β is unknown but we know LHα (corrected for attenuation) we can obtain the ultraviolet attenuation using equation (7.19), with a r.m.s. of 0.24 ± 0.02. This relation is valid under the assumption that the star formation rate is approximately constant in the last ≈ 108 yr. IIa) If we know Lobs Hα , but no estimate of A(Hα) is available, we can use Eq.(7.20) (rms∼ 0.34 ± 0.03). IIb) If neither β nor Hα luminosity are available we are left with the relations with H-band surface brightness6 (r.m.s.∼ 0.34 ± 0.03) and, in the worse case, III) H-band luminosity (rms∼ 0.39 ± 0.03 ). Summarizing, these relations allow us to estimate the value of the LT IR /LF U V ratio with an average uncertainties of∼0.32 dex. This value corresponds approximately to σ(A(F U V )) ∼0.5 mag, assuming log(LT IR /LF U V ) = 1 (the mean value for our sample) and using the model of Buat et al. (2005). This is the lowest uncertainty on the estimate of the LT IR /LF U V ratio in absence of far infrared observations. However we caution the reader that this value holds only for an optically-selected sample and that samples selected according to different criteria, especially FIR-selected, could contain higher dispersions. The cookbook presented in this chapter is obviously insufficient obtained for r.m.s. and m.a.d. are consistent with the ones expected for a Gaussian distribution 6 Since we need Hα flux to estimate metallicity, Eq.(7.15) cannot be used in this case. 7.6. A cookbook for determining LT IR /LF U V ratio x a 105 m.a.d.a b β 0.70 ± 0.06 1.30 ± 0.06 12 + log(O/H) 1.37 ± 0.24 −11.36 ± 2.11 LH /L 0.34 ± 0.10 −2.66 ± 0.88 µe (H) −0.28 ± 0.04 5.92 ± 0.81 LHα /LF U V 0.84 ± 0.07 −0.59 ± 0.12 obs LHα /LF U V 1.10 ± 0.17 −0.59 ± 0.21 rmsb 0.20 ± 0.02 0.26 ± 0.02 0.29 ± 0.03 0.25 ± 0.02 0.19 ± 0.02 0.27 ± 0.02 rs 0.26 ± 0.02 0.76 0.35 ± 0.03 0.59 0.39 ± 0.03 0.49 0.34 ± 0.03 −0.63 0.24 ± 0.02 0.76 0.34 ± 0.03 0.49 a: Mean absolute deviation from the best fit. b: Standard deviation from the best fit. Table 7.1: Linear realtions (log(LT IR /LF U V ) = a × x + b). useful to estimate the LT IR /LF U V ratio to understand dust attenuation and know how to correct UV observations of local and high redshifts galaxies, but it represents only the tip of the iceberg. The next steps should be the folowings: a) compare all the relations obtained in this work with different models in order to try to determine the physical properties of dust b) use models and data in order to estimate a new attenuation law from the far-ultraviolet to the near-infrared valid for normal star forming galaxies, as the one obtained for starbursts by Calzetti et al. (1994). Only knowing the dust attenuation law we will be able to correct for dust extinction all our observations and thus to correctly estimate the star formation rate in galaxies. Chapter 8 High velocity interaction: NGC4438 in the Virgo cluster This analysis represents the tip of the iceberg and only a future comparison with different dust models will allow us to understand dust attenuation and to know how to correct UV observations of local and high redshifts galaxies. A statistical analysis of star formation activity in cluster galaxies using UV data is therefore still impossible. For this reason, in the last three chapter of this work, I will focalize my attention on the study of three particular cluster galaxies considered as the prototypes of the three main environmental effects observed in clusters: tidal interaction, ram pressure stripping and preprocessing, respectively. These unique astrophysical laboratories will be used to deeply understand the effects of different physical mechanisms on galaxy evolution.. 8.1 Introduction NGC 4438 (Arp 120) is the clearest example of an ongoing tidal interaction in a nearby cluster of galaxies. Apparently located close to the Virgo cluster center (∼ 300 kpc from M87), NGC 4438 is a bulge-dominated late-type spiral showing long tidal tails (30 kpc) thought to be induced by a recent dynamical interaction with the nearby SB0 galaxy NGC 4435. Multifrequency observations covering the electromagnetic spectrum from X-rays (Kotanyi et al. 1983; Machacek et al. 2004) to radio continuum (Hummel & Saikia 1991), including both spectro-photometric and kinematical (Kenney et al. 1995; Chemin et al. 2005) data, have been carried out in the past to study the nature of this peculiar system. These observations have shown that the violent interaction between the two galaxies perturbed the atomic (Cayatte et al. 1990) and molecular (Combes et al. 1988) gas distribution, causing both gas infall toward the center which might have induced nuclear activity (Kenney et al. 107 108 8. High velocity interaction: NGC4438 in the Virgo cluster 1995; Kenney & Yale 2002; Machacek et al. 2004), and gas removal in the external parts displacing part of the gas in the ridge in between the two galaxies (Combes et al. 1988). Both multifrequency observational data (Kenney et al. 1995; Machacek et al. 2004) and model predictions (Combes et al. 1988; Vollmer et al. 2005) favor a recent (∼ 100 Myr) high-velocity, off-center collision between NGC 4435 and NGC 4438. Except for mild nuclear activity, it is still unclear whether the dynamical interaction between the two galaxies induced extra-nuclear star formation events: the low Hα/[NII] ratio and the similar X-ray and Hα morphology of NGC 4438 indicate that the Hα emission is in this case not due to the ionizing radiation but is probably due to gas cooling phenomena (Machacek et al. 2004). The UV emission is dominated by young stars of intermediate masses (2 < M < 5M ) and provides us with an alternative star formation tracer. As part of the Nearby Galaxy Survey (NGS), we have observed the central 12 deg2 of the Virgo cluster using the Galaxy Evolution Explorer (GALEX). A distance of 17 Mpc for Virgo is adopted. 8.2 Data The GALEX data used in this work include far-ultraviolet (FUV; λeff = 1530Å, ∆λ = 400Å) and near-ultraviolet (NUV; λeff = 2310Å, ∆λ = 1000Å) images. The data consist of 2 independent GALEX pointings centered at R.A.(J2000)= 12h29m01.2s, Dec(J2000)= +13◦ 10’29.6” (819 sec) and R.A.(J2000)= 12h25m25.2s, Dec(J2000)= 13◦ 10’29.6” (1511 sec), for a total of 2330 sec of integration time. To study the star formation history of NGC 4438, the UV data have been combined with visible and near-IR images of the galaxy taken from the GOLDmine database (Gavazzi et al. 2003a), from the SDSS Data release 3 (Abazajian et al. 2005) from the 2MASS survey (Jarrett et al. 2003) and from the CFHT and SUBARU archives. These are Hα+[NII] (Boselli & Gavazzi 2002), B (Boselli et al. 2003a), K’ (Boselli et al. 1997b), u, g, r, i, z SDSS, R CFHT and SUBARU and H 2MASS images. For the main body of the galaxy (region 4 in Fig. 8.1, see next sect.) we added the integrated spectrum (3500-7000 Å ; Gavazzi et al. 2004). The current calibration errors of the NUV and FUV magnitudes are on the order of ∼ 10% (Morrissey et al. 2005), comparable to that at other frequencies. 8.2. Data 109 Figure 8.1: The combined NUV and FUV image of NGC 4438. The regions described in sect. 3 of the text are labeled 1 to 7. The horizontal line is 10 kpc long (assuming a distance of 17 Mpc). 110 8.3 8. High velocity interaction: NGC4438 in the Virgo cluster The UV emission and the star formation history of NGC 4438 Figure 8.1 shows the UV image of NGC 4438, obtained by combining together the NUV and FUV frames in order to increase the S/N. The UV emission of the galaxy is mostly due to compact, bright regions in the central part of the galaxy (marked as region 4 in Fig. 8.1), in the northern tidal tail (region 2) and in the section of the southern tail closest to the main body of the galaxy (region 5). The UV emission is mostly diffuse in the extended western part of the galaxy (region 3) and at the edge of the southern tidal tail (region 6). In addition Figure 8.1 shows the presence of two extended and patchy emission to the north-west of the galaxy (∼ 15-25 kpc from the nucleus, marked as region 1 and region 7). These features, previously undetected in other visible and/or near-IR bands, are similar to a tidal tail: region 1 is ∼ 20 kpc long and ∼ 2 kpc wide, while region 7 is considerably smaller (∼ 2 kpc). The RGB image of the galaxy obtained by combining the FUV, NUV and B frames (see Fig.8.4) shows the color of the different regions: while the edge of both the northern and the southern tidal tails (region 3 and 6) are red (thus dominated by relatively old stars), regions 2 and 5 as well as the newly discovered regions 1 and 7, have blue colors and seem therefore to be dominated by a younger population. The Hα+[NII] emission map, given in Fig. 8.2 as a contour plot superposed on the NUV image of NGC 4438, shows a lack of massive, ionizing young O-B stars (Kennicutt 1998). The Hα+[NII] emission observed in region 5 has a different morphology than the UV one; on the contrary its distribution is the same observed in X-ray as stated by Machacek et al. (2004) (see Fig.8.3). This evidence confirms the conclusions of Machacek et al. (2004) that the Hα+[NII] emission is not due to the ionizing radiation but is probably associated with the cooling gas. What is the nature of the newly discovered extragalactic UV emitting regions? The average NUV surface brightness of these features is ∼ 28.5 ABmag arcsec −2 , while they are undetected both in the SUBARU R band (360 sec) image down to a surface brightness limit of ∼ 27.8 mag arcsec−2 and in Hα down to a surface brightness limit of ∼ 5 10−17 erg s−1 cm−2 arcsec−2 (see Fig.8.2), implying a log(N U V /Hα) ≥ 0.3. Extra-planar diffuse regions with an excess of UV over Hα flux ratio (i.e. log(N U V /Hα) ≥ 1, as that observed at 11 kpc from the disk of M82, are often interpreted as due to the UV radiation produced by the central starburst and locally scattered by diffuse dust (Hoopes et al. 2005). It is unlikely that scattered light is responsible for the UV emission since the steep slope of the UV spectrum (β=-2.32 and -2.05, as defined by Kong et al. 2004) is typical of a recent unreddened starburst (Calzetti 2001) and is unexpected in a scattering scenario since the dust albedo is greater in the NUV than in the FUV (Draine 2003) (i.e.β ≤ −1). Furthermore the lack of a powerful central starburst (as in M82) and the large distance of these 8.3. The UV emission and the star formation history of NGC 4438 111 Figure 8.2: The Hα+[NII] contours (red, in arbitrary scale, in between 8 10−17 and 6 10−16 erg cm−2 s−2 arcsec−2 , with σ= 5 10−17 erg cm−2 s−2 arcsec−2 , from Boselli & Gavazzi (2002)) are superposed to the NUV gray-level image of NGC 4438. Figure 8.3: Chandra image of NGC 4438 in the 0.3-2 keV energy band with Hα contours superposed. Adapted from Machacek et al. (2004) 112 8. High velocity interaction: NGC4438 in the Virgo cluster relatively patchy regions from the nucleus seem to exclude the scattering scenario. These data suggest that regions 1 and 7 are post starbursts, induced by the violent interaction with NGC 4435. In addition the absence of Hα emission associated with all the UV emitting regions suggests that the starburst lasted for a relatively short time, since it is not producing young, massive O-B stars any more. This is probably because the atomic and molecular gases, needed to feed star formation, have been removed during the interaction (Combes et al. 1988; Vollmer et al. 2005)1 . In order to date the starburst and reconstruct the star formation history of the galaxy, we have determined the spectral energy distribution (SED) of each region (see Fig.8.4) and then fitted it with a simple model of galaxy evolution. To this end we make the assumption that dust attenuating the SED is present only in region 4, where we correct the UV to near-IR data using the far-IR to UV flux ratio as done in Boselli et al. (2003a) and described in Appendix A. This restricted application is reasonable since no dust emission has been observed in the tidal tails with ISOCAM (Boselli et al. 2003b); furthermore, in regions 1 and 7, dust is unexpected since it has not yet been produced by the young stellar population, as confirmed by the steep β parameter (see also Chapter 7). Assuming that NGC 4438 was a normal late-type object before interacting with NGC 4435, we use the models of Boissier & Prantzos (2000) in order to reconstruct its SED before the interaction. The two parameters of the model (spin λ and rotational velocity VC ) are constrained by the observed total H-band luminosity and velocity rotation of NGC4438, leading to λ=0.01 and VC =290 km s −1 . In Fig.8.4 we compare the model with the SED of the main body of the galaxy (region 4), composed by an old population with no significant contribution from the recent starburst. Both the total SED and the optical spectrum produced by the model are in good agreement, confirming that the adopted technique is able to reproduce the galaxy SED before the interaction. We then assume that the evolved stellar population of each region, if present, is the one given by the model and removed from the main body of the galaxy by the tidal interaction, while the younger population is produced by the induced starburst. For each region, we thus combine the SED of an evolved stellar population with the one produced by an instantaneous burst of star formation obtained using Starburst 99 (Leitherer et al. 1999) for a solar metallicity and a Salpeter IMF between 1 and 100 M . For each age and intensity of the burst, we determine the best combination of evolved population+ burst by fitting the FUV to K band SED and rejecting solutions in disagreement (i.e. too bright) with the upper limits. We then adopt the age corresponding to the lowest reduced χ2 2 . The 1 The upper limit of the HI surface density for these regions is ∼ 1 M pc−2 (Cayatte et al. 1990) All ages with χ2 < 1 are acceptable solutions. Given the small number of photometric points available for regions 1 and 7 (2 GALEX bands), the fitted solution for a combination of a burst and an old population (two parameters) can be almost perfect (resulting in very low χ 2 , ≤ 10−2 ), as long as the obtained fit is in agreement with the limits at other wavelengths. Whenever the fit produces a SED not satisfying a detection limit, this solution is rejected. 2 8.4. Discussion and conclusion 113 results of our fitting procedure are presented in Fig.8.4. For each region (excluding region 4) two panels are given. The lower panels show the observed SED of each region (crosses, or arrows if are upper limits) and the best SED obtained from the fitting procedure (black line). The relative contribution of the evolved and young stellar populations to the observed SEDs are indicated in red and blue respectively. The burst luminosity contribution (for the age corresponding to the minimum χ2 ) in the band FUV, B and K is also given. In the upper panels the variation of the reduced χ2 parameter (black continuum line) and of the burst mass fraction (red dotted line) as a function of the age of the burst are given. This exercise gives an interesting result: the strong UV emission of regions 1 and 7 is due to a coeval starburst ∼ 6-20 Myr old. The age and the duration of the starburst are strongly constrained both by the lack of Hα emission and by the blue UV slope of the spectrum (lower limit to the age) and by the lack of an old stellar population (upper limit to the duration). The burst age for the other region cannot be determined with the same precision, but we can only put a lower limit to their age. Regions 2 and 5 are consistent with an older starburst (≥ 100 Myr, as suggested by their redder UV slope: β=-0.33 and -0.67 in regions 2 and 5 respectively) which probably ended ∼ 10 Myr ago as indicated by the lack of any Hα emission. Conversely the stellar population in regions 3 and 6 appear not significantly affected by the recent burst. Moreover it is interesting to note that, while the fraction of stars produced by this burst is dominant in regions 1 and 7, the sum of the stars produced by the burst in all regions (including the inner part) contributes to the total galaxy stellar mass by less than 0.1 %, an extremely low value for such a violent interaction. 8.4 Discussion and conclusion These observations have major consequences in constraining the evolution of cluster galaxies. A high-velocity off-center collision between two galaxies of relatively similar mass, whose violence is able to perturb the stellar distribution producing important tidal tails, is insufficient to induce a significant instantaneous starburst. This result might be representative only of the nearby Universe where encounters of gas-rich galaxies are probably rare since clusters are dominated by gas-poor early-type galaxies such as the companion galaxy NGC 4435. It is conceivable, however, that at higher redshifts, where clusters are forming, stellar masses produced by a starburst induced by interactions predicted by the models of Moore et al. (1996) (galaxy harassment) might be more important given the higher fraction of gas-rich galaxies. The other interesting result is the long time differential between the age of the interaction (∼ 100 Myr as determined by dynamical simulations, (Combes et al. 1988; Vollmer et al. 2005) and the beginning of the starburst (∼ 10 Myr in regions 1 and 7, ∼ 100 Myr in regions 2 and 5). This result is totally consistent with the models of 114 8. High velocity interaction: NGC4438 in the Virgo cluster Mihos et al. (1991) that predict for close-by encounters an enhancement of the star formation activity in the inner disk during some 100 Myr, stopping once the gas reservoir is exhausted as in NGC 4438. In the tidal tails, on the contrary, star formation is expected to increase after ∼ 100 Myr, the time needed by the gas to re-collapse, but then ceasing after a few Myr because the expansion of the tidal tail brings the gas surface density to subcritical values (no HI and CO has been detected in these regions). If these systems are dynamically stable and survive the interaction, they might be at the origin of some dwarf galaxies in the cluster similar to those observed in the Stephan’s Quintet by Mendes de Oliveira et al. (2004) or in other interacting systems (Neff et al. 2005; Hibbard et al. 2005; Saviane et al. 2004) Being produced by a single starburst, these gas poor systems might evolve into dwarf ellipticals, typical of rich clusters. Otherwise they will simply increase the fraction of unbound stars, contributing to the Virgo intracluster light (Willman et al. 2004). 8.4. Discussion and conclusion 115 Figure 8.4: The RGB (FUV=blue, NUV=green, B=red) color map of NGC 4438 and NGC 4435. The SED of each region defined in Fig. 1 are given in the lower plot of each frame. Crosses indicate the observed data, arrows upper limits (in mJy), the red dashed line the evolved population fit as determined by the model of Boissier & Prantzos (2000), the dotted blue line the starburst SED (from Starburst 99) and the dashed green line the combined fitting model. The burst luminosity contribution (for the age corresponding to the minimum χ2 ) in the band FUV, B and K is also given. The upper panel gives the variation of the reduced χ2 parameter (black continuum line, in logarithmic scale) and of the burst mass fraction (red dotted line) as a function of the age of the burst (in Myr). The lower panel of region 4 gives the integrated 3500 to 7000 Å, R=1000 spectrum of the main body of the galaxy (black continuum line) compared to the fitted model (red dashed line). Chapter 9 Ram Pressure stripping: NGC4569 in the Virgo cluster 9.1 Introduction Spiral disks can lose their atomic gas content during dynamical interactions with the hot and dense intergalactic medium (IGM) (Gunn & Gott 1972) and/or in direct interactions with nearby objects (Merritt 1983). These interactions can quench their star formation activity (Gavazzi et al. 2002c) leaving the objects to become anemic (van den Bergh 1976). To explain the well known morphological segregation effects Dressler (1980) it has been suggested that these quiescent spirals could evolve into lenticulars; however, observations and model predictions give still contradictory results (see Boselli & Gavazzi 2005 for a review). Despite the on-going physical processes (tidal interactions were probably dominant at early epochs, while galaxies-IGM interactions are more important at present), it is clear that the fuel supply needed to feed star formation is more efficiently removed where the host-galaxy potential well is weakest, i.e., in the outer disk. Given the strong relation between the gas surface density and the star formation activity in spiral disks, commonly known as the Schmidt law (Kennicutt 1998; Boissier et al. 2003), it is expected that star formation will be quenched in the outer (lower density) portions of the disk. While interferometric observations of galaxies in Virgo have clearly shown that HI disks are less extended in those objects located close to the cluster center (Cayatte et al. 1990), the observational evidence for a truncation of the star forming disks has been proven by Hα imaging (Koopmann & Kenney 2004b,a). Although a truncation of the disk profile has been predicted (Larson et al. 1980b), we still do not know what the passive evolution of a stellar disk is once its gas is removed. In particular, it is unclear whether the progressive radial suppression of star formation is able to reproduce the structural properties of lenticulars, generally characterized as 117 118 9. Ram Pressure stripping: NGC4569 in the Virgo cluster having higher surface brightness of their stellar disks and higher bulge-to-disk ratios than spirals (Dressler 1980). We have been collecting multi-frequency data for a large sample of late-type galaxies in nearby clusters and in the field in order to undertake comparative statistical analyses of any systematic differences between cluster and field objects. Combined with multi-zone models for the chemical and spectrophotometric evolution of galaxies (Boissier & Prantzos 2000), this unique database is helping us understanding the evolution of cluster spirals. As a first step during my thesis I studied the radial profiles of the Virgo cluster galaxy NGC 4569 (M90). NGC 4569, the prototype anemic galaxy as defined by van den Bergh (1976), is extremely deficient in HI, having just ∼ one tenth of the atomic gas of a comparable field galaxy of similar type and dimensions. The galaxy has a truncated Hα and HI radial profile (at a radius of ∼ 5 kpc; see Fig.9.3) as firstly noticed by Cayatte et al. (1994) and Koopmann & Kenney (2004a), witnessing a recent interaction with the cluster environment. NGC 4569 is located close (∼ 1 degree) to the cluster center. Being one of the largest galaxies (∼ 10 arcmin) in the Virgo cluster, NGC 4569 is the ideal candidate for our study since it can be spatially resolved at almost all wavelengths considered here. The combination of the multi-frequency 2-D data with our spectrophotometric models allow us to study, for the first time, the radial evolution of the different stellar populations in this prototype, gas-stripped cluster galaxy with the aim of understanding whether its structural properties can evolve into those of a typical cluster lenticular (S0) galaxy. 9.2 Data and models The large amount of spectrophotometric data available for NGC 4569, collected in the GOLDMine database (Gavazzi et al. 2003a), allow us to reconstruct its radial profile at different wavelengths: from the new GALEX UV bands (at FUV=1530 and NUV=2310 Å), to the visible B and V (Boselli et al. 2003a), Sloan u, g, r, i, z (Abazajian et al. 2005) and near-IR J, H and K bands (Boselli et al. 1997b; 2MASS Jarrett et al. 2003). Hα+[NII] narrow band imaging, used to trace the recent star formation activity, is available from Boselli & Gavazzi (2002). HI profiles are from Cayatte et al. (1994), while H2 profiles, determined from CO data using a luminosity dependent CO to H2 conversion factor (from Boselli et al. 2002b) are taken from the BIMA survey of Helfer et al. (2003) for the inner disk, and from Kenney & Young (1988) for the outer disk. The accuracy of the photometric imaging data is, on average ∼ 10 %. The galaxy rotation curve has been taken from Rubin et al. (1999). Unfortunately no metallicity gradient information is available for NGC 4569. The radial profiles have been constructed by integrating the available images within elliptical, concentric annuli. The ellipticity and position angles have been determined and then fixed using the deepest B band image following the procedure of Gavazzi 9.2. Data and models 119 et al. (2000) (see Fig.9.3). To avoid any possible contamination by the NW arm, whose kinematical properties indicate that it is not probably associated with the stellar disk but rather formed during the interaction with the ICM (Chung et al. 2005), the arm was masked in the construction of the radial profiles. If included, its contribution would be perceptible only in the FUV filter at radii > 8 kpc, increasing the surface brightness by < 0.5 mag. The UV to near-IR broadband images of the galaxy have been corrected for internal extinction using the recipe of Boissier et al. (2004), assuming a typical UV extinction gradient for a galaxy of the luminosity and scalelength of NGC 4569. In this case, in fact, the far-IR (IRAS) to UV flux ratio cannot be used to estimate the extinction because both fluxes are contaminated by the nuclear activity of the galaxy. The extinction in the other visible and near-IR bands has been determined using the prescription of Boselli et al. (2003a) and described in Appendix A. Hα+[NII] narrow-band imaging has been corrected for [NII] contamination and dust extinction (Balmer decrement) using the integrated spectroscopy of Gavazzi et al. (2004). To study the evolution of the disk of NGC 4569 at various radii, we have used the multi-zone chemo-spectrophotometric models of Boissier & Prantzos (2000), updated with an empirically-determined star formation law Boissier et al. (2003). These models have a resolution of ∼1 kpc, significantly lower than the one of our multiwavelength datasets (0.08-0.4 kpc). The errors in the surface brightness and color profiles have been computed following Gil de Paz & Madore (2005). For this reason we degraded all our images at the model resolution, and we extract the smoothed profiles used for the comparison between models and data. The nuclear emission due to the central AGN has been masked since the model is not able to reproduce the AGN activity (see Fig.9.4). The two model parameters (spin λ and rotational velocity VC ) are constrained by the H-band luminosity profile (determined assuming a distance of 17 Mpc) and the rotation curve of the galaxy, making the reasonable assumption that both of these observables are unperturbed during the interaction. This gives λ=0.04 and VC = 270 km s−1 (see Fig.9.1). To compute the Hα profile, the number of ionizing photons predicted by Version 5 of STARBURST 99 Vázquez & Leitherer (2005) for a single generation of stars distributed on the Kroupa et al. (1993) initial mass function (as used in our models) is convolved with our star formation history, and converted into a Hα flux as described in Appendix B. In addition to this model (valid for an unperturbed galaxy) we add an episode of ram pressure gas stripping. For simplicity, we adopt the plausible scenario of Vollmer et al. (2001): the galaxy has crossed the dense IGM only once, on an elliptical orbit. The ram pressure exerted by the IGM on the galaxy ISM varies with time (t) following a Lorentzian profile (see Fig.9.2): = 0 (∆t)2 ((∆t)2 + (t − t0 )2 ) (9.1) 120 9. Ram Pressure stripping: NGC4569 in the Virgo cluster Figure 9.1: The radial profile of observed (open symbols) and extinction-corrected (filled symbols) H-band surface brightness (left) and of the rotational velocity (center) used to constrain the model without interaction (represented by the black solid line). The total gas radial profile (right) predicted by the unperturbed model (solid black line) is compared to the observed one (green filled circles), obtained by summing the HI component (red line) to the molecular one (blue and light blue) and correcting for Helium contribution (× 1.4), and to the model including the interaction (black dashed line). where t0 is when the galaxy crosses the dense cluster core at high velocity and 0 is the value of ram pressure at t0 . Following Vollmer et al. 2001 we assume a width profile ∆t = 9 × 107 years. In order to determine the amount of stripped gas we make the hypothesis that the gas is removed at a rate that is directly proportional to the galaxy gas column density Σgas and inversely proportional to the potential of the galaxy, measured by the total (baryonic) local density Σpotential (provided by the Σgas . The two free model). The gas-loss rate adopted is then finally equal to Σpotential parameters in our model are then t0 and 0 . We make the further assumption that no extra star formation is induced during the interaction. 9.3 The star formation history of NGC 4569: model predictions Once the width of the interaction event, ∆t, is fixed, it is possible to choose simultaneously t0 and 0 because the amount of gas left and its radial distribution depend strongly on 0 while the resulting stellar light profiles depend mainly on t0 (see Fig.9.3 for some examples). If the cluster core crossing time is recent only the youngest stellar populations (emitting in Hα, whose age is ≤ 4 106 yrs, or far-UV, ≤ 108 yrs) have had time to feel the progressive radial suppression of the star formation activity. 9.4. Discussion and conclusion 121 Figure 9.2: Ram pressure stripping intensity (in arbitrary units) as a function of time (Eq.9.1). Adapted from Vollmer et al. (2001). Comparing model predictions with the spectrophotometric radial profiles of cluster galaxies can thus be used to date the dynamical interaction with the IGM. We thus fitted the data with models for different values of t0 and 0 . An important modification applied to the usual χ2 test is that its value was artificially put to 100 for any model predicting surface brightnesses in disagreement with observational limits (non detections at relatively large radii) in order to reject these solutions. The model best matching the properties of NGC 4569 (Fig.9.4) is characterized by 0 = 1.2 M kpc−2 yr−1 and t0 = 100 Myr. This is largely consistent with the dynamical models of Vollmer et al. (2004a), who obtained t0 ∼ 300 Myr. Although not reproducing perfectly the surface brightness profile, this model is able to qualitatively reproduce the truncation of the total gas disk profile (see Fig.9.1) and of the Hα and UV radial profiles (Fig.9.3) as well as the milder truncation observed at longer wavelengths. It is interesting to note that although older cluster core crossing epochs give more truncated disk profiles in the old stellar populations (B and i bands, blue dashed line), this is not the case in the gas profile which is modified by contributions from the recycled gas. 9.4 Discussion and conclusion The present work gives the first quantitative estimate of the structural evolution of stellar disks in cluster galaxies due to gas removal caused by a dynamical interaction 122 9. Ram Pressure stripping: NGC4569 in the Virgo cluster Figure 9.3: The radial profile of the observed (empty green circles) and extinctioncorrected (filled green circles) total gas, Hα, FUV (1530 Å), NUV (2310 Å), B and i surface brightness. The yellow shaded area marks the range in between the observed (bottom side) and extinction-corrected (top side) surface brightness profiles. Surface brightnesses are compared to the model predictions without interaction (black solid line) or with interaction for several 0 and t0 parameters. Equal maximum efficiency (0 =1.2 M kpc−2 yr−1 ) and different age: t0 =100 Myr, red continuum line (the adopted model); t0 =500 Myr, grey long dashed line, t0 =1.5 Gyr, dashed magenta line. Equal age (t0 =100 Myr) and different maximum efficiency: 0 =3 M kpc−2 yr−1 , blue dotted line; 0 =1/3 M kpc−2 yr−1 , orange dotted line. 9.4. Discussion and conclusion 123 of the galaxy with the IGM. Although the model only qualitatively reproduces the observed multi-wavelength radial profiles (the mismatch being attributed to resolution effects) it delivers a strong message concerning the passive stellar evolution of stripped disks. First of all it is clear that the truncation of the total gas disk profile is soon reflected in the young population stellar disk, confirming the predictions of Larson et al. (1980b). As observed in NGC 4569, this gas-stripped galaxy has a color gradient opposite to that of normal, isolated spirals, which generally have bluer colors in their outer disks (see Fig.9.5 and Fig. 9.4). NGC 4569 is bluest towards the center. The trend is especially true for colors tracing the relatively young populations (<∼ 108 yr); colors tracing populations older than the interaction event present the usual gradient (i.e., redder towards the center). The inversion of the color gradient, here observed for the first time in a cluster galaxy, is well reproduced by our model. The consequence of these findings in the interpretation of the evolution of cluster spiral galaxies is significant. One of the most intriguing and still open question regarding the effects of the environment on the evolution of galaxies is that of the origin of lenticulars, and their overabundance in the centers of rich clusters. Are lenticulars an independent population of galaxies formed in the primordial high-density environments, or were they spiral disks whose star formation activity has been quenched once their gas reservoir was removed by the unfavorable cluster environment? Although the second interpretation seems logical, simple statistical considerations in the seminal work of Dressler (1980) show that this idea is not supported by the observations: spirals have lower surface brightnesses and bulge-to-disk ratios than lenticulars, and spirals are rotationally supported while lenticulars are dynamically hotter systems. Structural (and kinematical) modifications must thus be invoked if spirals are to be transformed into lenticulars. The present work has shown for the first time how a galaxy-cluster IGM interaction is able to remove most of the gas reservoir inducing important structural modifications in the disk properties. We have in fact shown that, because of the differential radial stellar evolution of spiral disks, we expect that cluster spirals have (at least at short wavelengths) more truncated disk profiles, inverting the outer color gradient with respect to similar but unperturbed objects. The surface brightness of the disk, however, mildly decreases in Hα and in the UV bands while remains mostly constant at longer wavelengths even 5 Gyr after the interaction (Fig. 9.4d). The differential evolution of the stellar disk due to gas stripping alone is thus not able to reproduce the structural properties of present-day lenticulars. Gravitational perturbations, such as tidal interactions with other galaxies (Merritt 1983), interactions with the cluster potential well (Byrd & Valtonen 1990) or a mixture of both (called ‘galaxy harassment’ by Moore et al. 1996) must be invoked to reproduce the observed properties of nearby lenticulars. This new study and analysis is consistent with the idea that the present evolution of late-type galaxies in clusters differs from that at earlier epochs, where late-type galax- 124 9. Ram Pressure stripping: NGC4569 in the Virgo cluster ies were mostly perturbed by dynamical interactions (pre-processing and/or galaxy harassment; Dressler 2004, Moore et al. 1996) which were able to thicken the stellar disks thereby producing the present-day cluster lenticulars. We hope to confirm this original result in the near future once multi-frequency data come available for a statistical significant sample of late-type cluster galaxies. 9.4. Discussion and conclusion 125 (a) (b) (c) (d) Figure 9.4: The observed and model surface brightness (a), color (b) radial profiles of NGC 4569. In the model profiles the continuum lines are for models with gas removal, dashed lines for unperturbed models. c) the reduced χ2 as a function of the look-back time of the ram-pressure event for a few efficiencies 0 (M kpc−2 yr−1 ). Models were computed each 100 Myr for 0 < t0 < 500 Myr, 200 Myr for 1.5 < t0 < 0.5 Gyr and 1 Gyr for 6.5 < t0 < 1.5 Gyr; and each 0.2 M kpc−2 yr−1 efficiencies between 0.4 and 1.6 (only the more relevant are shown here). d) the variation of the effective surface brightness (mean surface brightness within Re , the radius containing half of the total light) and radius due to differential variation of the star formation history of NGC 4569. Open triangles are for the unperturbed model, the other symbols for different ages of the interaction (100 Myr, 1.5 and 5.5 Gyr). 126 9. Ram Pressure stripping: NGC4569 in the Virgo cluster Figure 9.5: The RGB (continuum subtracted Hα =blue, FUV=green, NUV=red) color map of NGC 4569 Chapter 10 Galaxy Pre-processing: the blue group infalling in Abell1367 10.1 Introduction In the previous two chapters we have investigated the effects of the environment on the properties of galaxies inhabiting the core of the Virgo cluster. However galaxies interact with the harsh environment well before having reached the center of a cluster. In particular, if we believe that structures grow hierarchically, galaxy clusters form not by accreting individual galaxies randomly from the field, but rather through the infall of less massive groups falling in, along large scale filaments. Galaxy groups may therefore represent a natural site for a preprocessing stage in the evolution of cluster galaxies. These infalling groups have velocity dispersions that are significantly smaller than that of cluster, permitting the slow gravitational interaction typically observed in field galaxies. Moreover even in compact groups ram pressure seems to be able to displace the gas from the disk of galaxies (Fujita 2004; Roediger & Hensler 2005). This means that probably at least part of the morphological and star formation properties of cluster galaxies derives from earlier epochs and very different conditions than the ones observed in today clusters (Dressler 2004). Environmental interactions in the infalling groups may thus represent a preprocessing step in the evolution of cluster galaxies (Mihos 2004a). Unfortunately, witnessing preprocessing in local Universe is a real challenge since we live in a Λ-dominated Universe where the infall rate is significantly lower than in the past (Gottlöber et al. 2001). Today, we observe a plethora of clusters experiencing multiple merging (Gavazzi et al. 1999a; Donnelly et al. 2001; Cortese et al. 2004), but the structures involved are subclusters with a mass ∼ 5 × 1014 M , considerably higher than the typical mass of a compact group ∼ 1013 M (Mulchaey 2000), as the North and South subclusters in Abell1367 studied in Chapter 5 (see Table 5.7). However Abell1367 represents a unique excep127 128 10. Galaxy Pre-processing: the blue group infalling in Abell1367 tion among local, dynamically young, clusters since in addition to massive evolved substructures it is also experiencing the merging of a compact group infalling directly into the cluster core. This group has a velocity dispersion of only ∼ 170km s−1 , and it is infalling into the cluster core at a very high speed (∼ 1700km s−1 ). The rarity of this phenomenon could probably explain the unique properties observed in this group. In fact it was independently discovered by Iglesias-Páramo et al. (2002) and Sakai et al. (2002) during two deep Hα surveys of nearby clusters, representing the region with the highest density of star forming systems ever observed in the local Universe. Sakai et al. (2002) argued that this group lies in the cluster background, having no interaction with the cluster environment. On the contrary the dynamical analysis presented in Chapter 5, is consistent with an infalling scenario, as also proposed by Gavazzi et al. (2003b). Moreover this picture is supported by X-ray observations: Sun & Murray (2002) (using Chandra observations) discovered extended gas features and a ridge near the SE cluster center. They proposed that these features are associated with a new merging component penetrating the SE subcluster. XMM clearly detects a cold front near the center of the SE subcluster, probably associated with a group infalling into the cluster core (A. Finoguenov, private comm.). All these observational evidences suggest that we are witnessing, for the first time in the local Universe, a compact group infalling into a core of a dynamically young cluster. It thus represents a unique laboratory to study with the great detail possible only in the local Universe, a physical process typically expected in clusters at high redshift. The study of this group could therefore help us shading light on the possible influence that preprocessing might have and have had on the past evolution of galaxies now populating high density environments. During the last few years we thus collected a great amount of multiwavelength spectroscopic and imaging observations in order to try to reconstruct the history of this rare group of galaxies, which represents the only compact group infalling into the center of a galaxy cluster ever observed in the nearby Universe. Throughout this chapter I will refer to this group as the Blue Infalling Group (BIG), as defined by Gavazzi et al. (2003b) 10.2 Observations 10.2.1 HI observations Using the refurbished 305-m Arecibo Gregorian radio telescope we observed the BIG region in March 2005. We obtained observations for 4 different positions covering the group center and its NW outskirt (see Fig.10.1). Data were taken with the L-Band Wide receiver, using nine-level sampling with two of the 2048 lag subcorrelators set to each polarization channel. All observations were taken using the position-switching 10.2. Observations 129 Figure 10.1: The four Arecibo HI pointings obtained in the region of the BIG group, superposed to the r 0 band image. The size of each circle correspond to the telescope beam. 130 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Figure 10.2: GALEX NUV image of the Blue Infalling group (BIG). technique, with each blank sky (or OFF) position observed for the same duration, and over the same portion of the telescope dish as the on-source (ON) observation. Each 5min+5min ON+OFF pair was followed by a 10s ON+OFF observation of a wellcalibrated noise diode. The velocity resolution was 2.6 km s−1 , the instrument’s beam at 21 cm is 30.5×30.1 and the pointing accuracy is about 1500 . Flux density calibration corrections are good to within 10% (and often much better), see the discussion of the errors given in O’Neil (2004). Using standard IDL data reduction software available at Arecibo, corrections were applied for the variations in the gain and system temperature with zenith angle and azimuth. A baseline of order one to three was fitted to the data, excluding those velocity ranges with HI line emission or radio frequency interference (RFI). The velocities were corrected to the heliocentric system, using the optical convention, and the polarizations were averaged. All data were boxcar smoothed to a velocity resolution of 12.9 km s−1 for further analysis. 10.2. Observations 10.2.2 131 UV to near-IR imaging The Blue Infalling Group has been observed by GALEX in April 2004, within the two pointings of the Abell cluster 1367. The observations are centered at R.A.(J2000)= 11:43:41.34 Dec(J.2000)=+20:11:24.0 (e.g. offset to the north of the cluster to avoid a star bright enough to threaten the detector), with a mean exposure time of 1460s, as described in Chapter 4. Fig.10.2 shows the GALEX NUV image of the Blue Infalling Group. UBVRH photometry for CGCG (Zwicky et al. 1961) galaxies is taken from Gavazzi et al. (2003a). 10.2.3 Hα imaging We observed BIG using the Device Optimized for the LOw RESolution (DOLORES) attached at the Nasmyth B focus of the 3.6m TNG in the photometric nights of 17th May and 18th June, 2004. The observations were taken through a [SII] narrow band filter centered at ∼ 6724Å and a width of ∼ 57Å covering the redshifted Hα and [NII] lines. The underlying continuum was taken through a broadband (Gunn) r 0 filter. Images, split in 6 exposures of 1200 sec in the narrow band filter and 5 exposures of 300 sec in the r 0 broadband filter, for a total of 2 hours and 30 minutes exposure respectively, were taken with a seeing of ∼ 1.2 arcsec. The photometric calibration was achieved by exposing the spectrophotometric star Feige 34. After bias subtraction and flat-fielding, the images were combined. The intensity in the combined OFF-band frame was normalized to that of the combined ON-band one by the flux ratio of several field star. The NET image was obtained by subtracting the normalized OFF-band frame to the ON-band one. The resulting OFF and NET-band frames are shown in Figs. 10.6 and 10.7 respectively. Hα+[NII] fluxes and EWs are obtained as described in Boselli et al. (2002a). 10.2.4 MOS spectroscopy We observed the BIG region in MOS mode with the ESO/3.6m and with the TNG telescope. The ESO/3.6m observations were taken in the photometric nights of May 5th and 6th 2003 with the ESO Faint Object Spectrograph and Camera (EFOSC). We used the MOS mode of EFOSC to obtain the spectra of 9 of the emitting line knots. The EFOSC spectrograph was used with a 300 gr/mm grating and 2048×2048 thinned Loral CCD detector, which provided coverage of the spectral region 3860 − 8070 Å. Slits width of 1.75” yielded a resolution of ∼ 19Å. We obtained eleven exposure of 1530 sec, for a total exposure time of ∼ 4.65 hours. The TNG observations were taken in the photometric nights of 26th March and 22nd April 2004 with DOLORES. We used the MOS mode of DOLORES to obtain the spectra of 8 of the emitting line knots, of the nuclear region of CGCG97-125 and of 14 132 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Figure 10.3: High-contrast Hα+[NII] band frame of the BIG group. 10.2. Observations Name R.A. (J2000) 133 Dec (J2000) TNG K1 DW3 d DW3 e DW3 a 97-114b 97-114a K2 a K2 b DW2 c DW2 b DW2 a K5 DW1 b DW1 c DW1 a 97-125b K3 97-125a 114444.18 114445.97 114445.97 114446.43 114446.56 114447.41 114450.61 114449.71 114451.12 114451.17 114451.67 114451.76 114453.78 114454.29 114454.64 114454.89 114455.28 114455.99 194816.0 194744.4 194741.1 194741.2 194640.3 194649.8 194605.1 194604.7 194718.7 194717.5 194713.5 194752.7 194731.5 194728.6 194732.9 194611.3 194803.3 194628.0 8422 ± 153 − − 8490 ± 180 − − − 8309 ± 165 8380 ± 188 − 8253 ± 292 − − 8343 ± 223 − 8261 ± 191 8020 ± 212 − Velocity (km s−1 ) ESO − MOS Sakai02 Gavazzi03 8265 ± 117 8564 ± 151 8072 ± 124 − 8656 ± 132 8763 ± 124 8080 ± 140 − − 8221 ± 146 − 8241 ± 112 − − 8265 ± 136 8396 ± 132 − − − − − 8266 8504 − 8070 − − − − − 8070 − 8161 8170 − − 8098 − − − 8383 8425 8089 − − 8077 − 7995 − − 8067 − − 8330 Table 10.1: Redshifts of the galaxies in the BIG group. galaxies in the region. The DOLORES spectrograph was used with a grating which provided coverage of the spectral region 3200 − 8000Å. Slits width of 1.6” yielded a resolution of ∼ 17Å. We obtained six exposure of 1800 sec, for a total exposure time of 3 hours. All the emitting line regions observed in MOS spectroscopy are shown in Fig.10.3 In addition we took spectra of the bright galaxy CGCG97-114 using the Loiano/1.52 m telescope. The BFOSC spectrograph attached at the Loiano telescope was used with a 300 gr/mm grating and 1300 × 1340 thinned EEV CCD detector, which provided a spectral coverage 3600 − 8900Å. A slit width of 2.00” yielded a resolution of ∼ 20Å. The observations were taken in the ”drift-scan” mode, with the slit parallel to the galaxy major axis, drifting over the optical surface of the galaxy. The total exposure time was 2400 sec. The reduction of the spectra was carried out using standard tasks in the IRAF package. Bias subtraction and flat-field normalization was applied using median of several bias frames and flat-field exposures. The various exposures were combined using a 134 10. Galaxy Pre-processing: the blue group infalling in Abell1367 median filter, thus removing the cosmic rays. The λ calibration was carried out using IDEN T IF Y −REIDEN T IF Y −F IT COOR on exposures of He/Ar lamps for each slit, and the calibration was transferred to the science frames using T RAN SF ORM . Typical errors on the dispersion solution are of ∼ 0.5 − 1Å, as confirmed by the measurements of the sky lines. However, since the resolution of our spectra is ≥ 13Å we assume an rms of 3Å on our wavelengths calibration. The two-dimensional frames were sky subtracted using BACKGROU N D. One-dimensional spectra were obtained integrating the signal along the slit using AP SU M . The apertures were limited to regions where the signal intensity was above 1 σ of the sky noise. Spectra were flux-calibrated using the spectrophotometric standard star: ltt 3864 for the ESO, Feige 67 for the TNG and Feige 34 for the Loiano observations. The redshift of each knot was derived as the mean of the individual redshift obtained from each emission line. Our results are shown in Tab. 10.1 and compared with the previous measurements by Sakai et al. (2002) and Gavazzi et al. (2003b). Line measurements All spectra were shifted to the rest frame wavelength and normalized to their intensity in the interval 5400-5600 Å. The flux-calibrated, normalized spectra are presented in Fig. 10.18. Under visual inspection of the spectra we carried out the measurement of the emission lines using SP LOT . This provided a list of fluxes and EWs with respect to a user defined continuum level. Hα (λ6563) is bracketed by the weaker [NII] doublet ([NII1] λ6548 and [NII2] λ6584). The three lines are not well resolved, thus using the task SPLOT we performed a two Gaussian fit to the blended emissions providing an estimate of the line ratio [NII]λ6584/(Hα + [NII]λ6548). The two bright galaxies CGCG97-125 and CGCG97-114 show evidence for underlying absorption in correspondence to emission lines. We de-blended the underlying absorption from the emission lines as discussed in Chapter 7. In order to compare our observations with the ones presented by Sakai et al. (2002) we re-measured, using the method described above, the spectra taken at the Stewart Observatory 2.3m Bok telescope and at the 6.5m MMT by these authors. The two sets of measurements presented in Tab. 10.2 are found in fair agreement. 10.2.5 High Resolution spectroscopy We obtained high dispersion long-slit spectra of CGCG97-125 and CGCG97-120 with the 1.93 mtelescope of the Observatoire de Haute Provence (OHP), equipped with the CARELEC spectrograph coupled with a 2048×512 TK CCD, giving a spatial scale of 0.54 arcsec per pixel. The observations were carried out in the night of April 20, 2004 in approximately 2 arcsec seeing conditions through a slit of 5 arcmin × 2 arcsec. The selected grism gives a spectral resolution of 33 Å/mm or 0.45 Å/pix and a spectral 10.3. Results Object K1 K1 DW3 a DW3 d 97-114b 97-114a 97-114 K2 a K2 b DW2 a DW2 b K5 DW1 b DW1 c DW1 a DW1 a 97-125b 97-125b 97-125 97-125 K3 135 Tel. ESO T NG T NG ESO ESO ESO LOI ESO T NG T NG ESO ESO MMT T NG ESO MMT ESO T NG OHP T NG T NG C1 [OII] (0.00) 3.68 0.00 3.49 0.02 8.26 0.00 2.01 0.33 4.25 0.24 3.37 0.75 2.61 0.17 4.06 0.23 8.43 0.16 5.45 > 0.1 > 5.18 0.56 − 0.33 3.62 0.00 2.76 0.30 − 0.20 3.75 0.55 − 0.04 3.60 0.88 9.23 0.90 6.58 0.00 2.87 Hβ [OIII] [OIII2] Hα [NII2] 1.00 0.94 1.00 0.95 1.00 0.31 1.00 0.60 1.00 0.57 1.00 0.26 1.00 0.21 1.00 0.77 1.00 0.64 1.00 0.73 1.00 < 1.00 1.00 0.47 1.00 0.74 1.00 0.47 1.00 0.82 1.00 0.80 1.00 0.51 1.00 0.35 1.00 1.06 1.00 1.15 1.00 0.25 2.53 2.57 1.06 1.34 1.83 0.70 0.36 2.09 0.82 0.99 < 0.99 0.65 2.50 1.53 2.39 2.41 1.25 1.28 1.97 1.89 0.97 − − 2.86 0.58 2.86 0.59 2.86 0.69 2.86 0.26 2.86 0.45 2.86 0.64 2.86 0.50 2.86 0.32 2.86 0.60 2.86 < 0.61 2.86 0.29 2.86 0.35 2.86 0.66 2.86 0.35 2.86 0.49 2.86 0.37 2.86 0.60 2.86 0.84 2.86 1.21 2.86 0.89 Table 10.2: Line fluxes, corrected for internal extinction, of the galaxies in the BIG group. coverage in the region 6080-6990 Å containing the redshifted Hα ( λ 6562.8 Å), the [NII] doublet (λ 6548.1, 6583.4 Å) and the [SII] doublet (λ 6717.0, 6731.3 Å). 10.3 Results 10.3.1 Kinematics Table 10.1 lists the positions and radial velocities of the objects that were measured spectroscopically. Our observations confirm the physical association of all the emitting line objects with the bright galaxies CGCG97-114 and CGCG97-125. On the contrary the brightest galaxy in this region CGCG97-120 seems not associated with this group, having a recessional velocity of 5635 km s−1 (see also Section 10.3.7). The velocity of galaxies in BIG (< V >= 8230 km s−1 and σV = 170 km s−1 ) exceeds significantly the mean cluster velocity of < V >= 6484 km s−1 (σV = 891 km s−1 ), 136 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Figure 10.4: Upper panel: The position and the width (rectangular areas on the right) of the three slits obtained for CGCG97-125. The slits are superposed to the Hα + [NII] net image. Lower Panel: The three different rotations curves obtained for CGCG97-125. Letters indicate the different regions as labeled in the upper panel. 10.3. Results 137 Figure 10.5: The low resolution 2D spectrum obtained at ESO/3.6 for the knots DW3d (left) and DW3e (right), shows a significant difference (∼ 500km s−1 ) in the velocity of the two knots. suggesting that it is infalling at ∼ 1700 km s−1 into the cluster core. The high resolution spectra obtained at the OHP telescope give us more insights on the dynamical state of CGCG97-125. Velocity plots of CGCG97-125 were extracted from each spectrum by measuring the wavelength of the Hα line in each pixel along the slits. The three rotation curves so obtained are given in Fig.10.4. In each diagram the recessional velocity is plotted as a function of position along the slit (the spatial axis runs from E (left) to W (right)). All the three spectra show regions with multiple velocity components, especially in correspondence to the galaxy center where two sudden velocity jumps of ∼100-150km s−1 are clearly present. It is interesting to note that the velocity of these jumps decrease from ∼ 8400 km s−1 to ∼ 8150 km s−1 and their position moves to east, passing from the north to the south part of the galaxy. Even if several examples of kinematic disturbances has been observed in normal galaxies (Rubin et al. 1999; Haynes et al. 2000) and interacting systems (Jore et al. 1996; Duc & Mirabel 1998), the features observed in CGCG97-125 are extremely rare. To our knowledge, the only other galaxy with the same characteristics is UGC6697 (Gavazzi et al. 2001b), the merging systems in the NW part of Abell1367 (see Chapter 5). The velocity jumps observed in the rotation curve of CGCG97-125 are consistent with the idea that this galaxy has experienced a merging in the past; however its properties are unusual if compared with what expected from a similar phenomenon. During the accretion of a satellite, the gas falling into the galaxy center is expected to relax before the gas at the outskirts of the galaxy. The relaxation time is in fact ∝ R/V, where R is the radial distance from the center and V is the rotational velocity. On the contrary in this case, the major anomalies are observed near the galaxy center while in the outer part the rotation curve presents a 138 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Name r0 mag Hα flux erg cm−2 s−1 97125 97114 DW1 DW2 DW3 13.99 15.06 17.93 18.96 19.11 (1.33 ± 0.29) × 10−13 (6.59 ± 0.71) × 10−14 (1.60 ± 0.17) × 10−14 (3.67 ± 0.86) × 10−15 (4.47 ± 0.96) × 10−15 EW (Hα + [NII]) SF R a Å M yr−1 27 ± 3 34 ± 5 128 ± 15 25 ± 7 56 ± 15 1.49 0.74 0.09 0.02 0.02 a: obtained using equation 7.7. L(Hα) corrected for [NII] contribution and extinction using values obtained from spectroscopy (see Table 10.2). Table 10.3: Properties of galaxies in BIG. typical S shape. Detailed dynamical simulations of a minor merger experienced by an S0 galaxy are thus mandatory to try to understand the particular features observed in this galaxy. The MOS spectroscopy collected at the ESO/3.6m and at the TNG telescopes gives us some information regarding the internal dynamic of DW3. The emitting line knots composing this system have considerably different recessional velocities, ranging between ∼ 8000 km s−1 and ∼ 8600 km s−1 . The western (DW3-b) and the eastern (DW3-a) knots have a recessional velocity of ∼ 8250 − 8300 km s −1 significantly lower than the one observed in the northern knot DW3-d (∼ 8564 km s−1 ) and ∼ 250 km s−1 higher than the redshift of the southern knot DW3-e. This great difference is clearly visible in Fig.10.5 where the two emitting line knots DW3-d and DW3-e are observed within the same slit (thus the relative offset is not affected by any uncertainty in the wavelength calibration). The observed high velocity gradient (∼ 500 km s−1 ) suggests that these five knots are probably not gravitationally bound, and thus that DW3 does not represents a dwarf virialized system. 10.3.2 Hα properties When observed in optical broad band images this group does not show any unexpected feature if compared with other group of galaxies. The r 0 -band luminosity function of BIG in the interval -19.2 < Mr0 < -12.2 has a slope α ∼ -1, consistent with the r 0 -band luminosity function of Hickson compact groups (Hunsberger et al. 1998), suggesting that originally BIG was a normal compact group. On the opposite, BIG represents a true exception as far as its Hα properties. At least ten out of 12 star-forming regions are associated with dwarf systems (or extragalactic HII regions) with E.W.(Hα + [NII]) often exceeding 100 Å(Gavazzi et al. 2003b). These galaxies, in spite of being ∼1000 times smaller than typical giant galaxies, are currently forming stars at a 10 times higher rate (per unit mass) than normal galaxies of similar 10.3. Results 139 Figure 10.6: Stellar shells are seen around galaxy 97-125 in the r 0 band image of BIG. No continuum emission is detected from the low brightness trails (except K2). 140 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Figure 10.7: Extended low brightness trails appear in the Hα+[NII] NET frame of BIG. 10.3. Results 141 luminosity, as derived from their L(Hα) (see Table 10.3). As remarked by Sakai et al. (2002), it is the first time that such a high density of star-forming galaxies has been seen in a nearby cluster, in spite of having collected data over an area of A1367, Coma, and the Virgo Cluster approximately 500 times larger than the group size. Moreover new Hα images of BIG obtained last year reinforce the uniqueness of this group revealing a spectacular Hα filamentary structure on top of which the star forming knots observed by Sakai et al. (2002) and Gavazzi et al. (2003b) represent the tip of the iceberg. Multiple loops of ionized gas appear with a projected length exceeding 150 kpc, a typical transverse size of 5 kpc, among the most extended low-brightness Hα emission features ever detected (see Fig. 10.7). One stream (labeled NW in Fig. 10.7) extends from the northern edge of the frame to the dwarf galaxy DW3, with an extension of ∼ 100kpc. The second and brightest one (labeled W in Fig. 10.7) traces a loop around galaxy 97-120 and seems connected to the bridge (labeled K2 in Fig.10.7) between 97-114 and 97-125, that was known from previous studies. If this is the case, the total projected extension of the NW and W trails would result ∼ 150 kpc. In addition to the filamentary features, at least two other diffuse Hα regions (labeled S and E in Fig.10.7) are detected. The total diffuse (Hα + [NII]) emission (e.g. excluding the contribution of the three bright galaxies and of the ten dwarfs/HII regions previously discovered) results ∼ 1.2×10−13 erg cm2 s−1 i.e. similar to the flux collected from one of the bright galaxies, and the typical surface brightness is 10−17.6 − 10−18.3 erg cm−2 s−1 arcsec−2 . Along the filaments we detect typically an E.W.(Hα + [NII]) ≥ 100 − 150 Å. The loop around 97-120 alone contributes with ∼ 2.4 × 10−14 erg cm2 s−1 , as obtained integrating the Hα + [NII] emission in a circular corona of 10 kpc radius and an annulus of 5 kpc, centered on 97-120. The derived line intensity is 2.05 Rayleigh (1 Rayleigh = 106 /4π photons cm−2 s−1 sr−1 ), corresponding to an emission measure (EM) of 5.7 cm−6 pc. Assuming a torus geometry with a circular section of radius ∼ 5 kpc and a filling factor of 1, the plasma density results ne ∼ 3.3 × 10−2 cm−3 and the ionized column density Ne ∼ 5×1020 cm−2 (the inferred densities would be higher if the gas is in clumps or filaments, which is likely). The emission measure in the NW trails results lower (∼ 1.3 cm−6 pc) than in the loop around 97-120 and the plasma density is ne ∼ 1.1 cm−3 . The trails geometry is strongly suggestive of a rosetta orbit typical of tidal disruption of a satellite galaxy. However contrary to other known examples of tidal streams the features here observed show strong Hα emission and no continuum emission above Σr0 = 26.8 mag arcsec−2 (even though this limit is insufficient to rule out the presence of stellar streams of brightness as low as observed for example in the M31 stream (Ibata et al. 2001)). The case offered by BIG seems therefore unique as it combines the, eventually present, faint stellar brightness of tidal streams with strong line emissions of tidal tails. What mechanism have produced a such unusual feature? In order to test the tidal disruption scenario we use the formalism of Johnston et al. (2001), 142 10. Galaxy Pre-processing: the blue group infalling in Abell1367 assuming: 1) that at least one of the gas trails is from a dwarf intruder merged into 97-125 and 2) that the geometry of the undetected stellar streams is the same of the observed gaseous trails. The intrinsic geometry of a streamer from a totally disrupted satellite can be used to estimate the mass m and age t of a young streamer: 2 3 R v p circ 11 w M , (10.1) m ∼ 10 R 10 kpc 200 km/s and the time since its disruption R Rcirc 200 km/s t ∼ 0.01 Ψ Gyr , w 10 kpc vcirc (10.2) where w is the width of the streamer at radius R, Ψ is its angular length, Rp is the pericentric distance of the orbit and Rcirc is the radius of the circular orbit with the same energy as the true orbit. Of course, we cannot measure Rcirc directly, but we can approximate it as being halfway between the adopted apocenter and pericenter. Thus adopting a projected ratio of the loop width w, to the radius R, of ∼ 0.15, a pericentric distance of Rp ∼ 15 kpc, an orbit with the same energy of a circular orbit of radius Rcirc ∼ 30 kpc and a rotation velocity Vcirc ∼ 298km s−1 (Vogt et al. 2004), we obtain a satellite mass ∼ 1 × 109 M and an age of the interaction ∼ 1.3 Gyr. The mean surface brightness of the tidal debris is then obtained using the following equation: 1 Gyr 10M /L,ν µr0 (t) = −2.5 log f Υ t 2.5 m 10 kpc vcirc + 23.9 + M,r0 ,(10.3) − log 3 200 km/s 108 M R where M,r0 is the r 0 absolute magnitude of the Sun, Υ is the r 0 mass to light ratio of the satellite and f is the mass fraction loses by the satellite. The mean surface brightness of the tidal debris in BIG is f 3 (10.4) µr0 = 26.8 − 2.5 log Υ 1 consistent with the undetection of continuum emission above Σr0 = 26.8 mag arcsec−2 . However I stress the reader that this simulation is based on the interaction between two field galaxies and not infalling into the cluster center as in this case. The most dramatic difference between field mergers and those in a cluster is in the evolution of tidal debris. In the field, most of the material stripped into tidal tails remains loosely bound to the host galaxy, forming a clear tracer of the gravitational interac- 10.3. Results 143 tion. In the cluster encounter, the cluster tidal field quickly strips the material from the galaxy, dispersing it throughout the cluster and making these tidal tracers very short-lived (Mihos 2004a). Thus we can assume the obtained value of µr as a lower limit for the surface brightness in the stellar trails. Although the undetection of stellar emission in the trails does not help us ruling out a tidal stream nature for these trails, their strong Hα emission makes BIG a unique example among interacting systems and compact groups. Conversely other known tidal tails have E.W.(Hα) ranging from zero (i.e. the Stephan’s quintet) to ∼ 20 Å (i.e. the Mice (NGC 4676) and some Hickson compact groups). Moreover tidal streams discovered in interacting systems (e.g. Shang et al. 1998; Forbes et al. 2003) are detected only in continuum with no Hα emission, even if associated to strong starburst merging systems (Wehner & Gallagher 2005). For these reasons, the unique features observed in BIG make us suppose that not only tidal interaction can produce the Hα trails but that probably the mutual influence of tidal and non-gravitational forces (e.g. ram pressure) can explain the physical properties of this group. In order to explain the properties of these trails we need a mechanism able to strip gas from galaxies with little or no influence on the stellar component: a condition respected only by galaxy interaction with the hot intracluster medium. This scenario is also supported by the discovery in the NW part of Abell1367 of two low surface brightness Hα cometary tails, with a total length of 75 kpc, associated with two star forming systems: CGCG 97-073 and 97-079 (Gavazzi et al. 2001a). In fact, the morphology and properties of the tails (which typical size and gas densities are similar to the trails observed in BIG) suggests that galaxies in the NW group are experiencing ram pressure due to their high velocity motion through the IGM. The only difference between these two cases is that CGCG97-079 and CGCG97-073 are infalling into the cluster as isolated systems, while galaxies in BIG are infalling within a compact group where gravitational interactions are not negligible. 10.3.3 HI properties HI observations give us additional hints on the properties of this unique group. CGCG97-125, the brightest member of BIG, has a normal hydrogen content: M(HI) = 3.9 × 109 M (Sakai et al. 2002), implying an HI deficiency1 = -0.21. Its HI column density distribution appears asymmetric, with the highest signal in the western side of the galaxy, as seen in the HI map obtained by Sakai et al. (2002) and reproduced in Fig.10.8, suggesting that this galaxy is strongly perturbed by an external agent. On the contrary CGCG97-114 has an HI mass of only M(HI) = 3.0 × 108 M (Sakai et al. 2002) with a resulting HI deficiency of 0.7. This low content is surprising if 1 The HI deficiency is defined as the difference, in logarithmic units, between the observed HI mass and the value expected from an isolated galaxy with the same morphological type T and optical obs obs linear diameter D: HI DEF = < log MHI (T obs , Dopt ) > −logMHI (Haynes & Giovanelli 1984) 144 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Figure 10.8: HI column density distribution in BIG. Contours are 0.5, 1.0, 2.0, 3.0, 4.0, 5.0, and 6.020 cm−2 . Adapted from Sakai et al. (2002) Figure 10.9: HI position-velocity diagram centered on CGCG 97-125. Adapted from Sakai et al. (2002) 10.3. Results 145 compared with the high star formation activity (0.74 M yr−1 , see Table 10.3). At this current SFR, the total HI mass of CGCG 97-114 would be depleted in 2.5 × 108 yr or in 9.4 × 108 if we add the total molecular gas mass (4 × 108 ) detected by Boselli et al. (1997a) in this galaxy. This suggests that the galaxy is currently experiencing an intense, transient burst of star formation. In addition to the two detected CGCG galaxies, in Fig.10.8 there appears to be extended HI, mostly around CGCG 97-125. The HI extension appears to be a continuation of the Hα structure to the west of CGCG 97-125. This extended structure is typical of galactic merger remnants (Hibbard & van Gorkom 1996) and suggests that a recent merger has affected this galaxy. The HI distribution around CGCG 97-125 is extended not only in the plane of the sky but in the velocity dimension. The position-velocity diagram presented by Sakai et al. (2002) centered on CGCG 97-125 is shown in Figure 10.9. The velocity distribution shows a regular gradient across the galaxy (the optical major axis of CGCG 97-125 is very close to east-west) ranging from 8090 up to 8490 kms−1 , corresponding to a rotation speed of 298 km s−1 when corrected for inclination (Vogt et al. 2004). This value is exceptionally high for a galaxy of the same luminosity, which usually has a rotation speed of ∼200 km s−1 . Thus, both the HI distribution and the HI kinematics yet available suggest that CGCG 97-125 is a quite peculiar object. Addition information concerning the HI properties of this group can be obtained from Arecibo observations. In Fig.10.10 are shown the four spectra obtained for the different pointings of BIG. Unfortunately three of the four pointings (97-125, 97-120, 97-114) are surely not independent due to the large overlap in the observed fields. The only, if any, independent observation is represented by the NW field, since it is far away from all bright galaxies and has a relatively small overlap with the field centered on 97-120. Since the side-lobes are located ∼5 arcmin from the field center, we can exclude a strong contamination of the NW pointing from the bright galaxies in BIG. In addition to the strong HI emission in the velocity range 8000−8500km s−1 , associated to the star forming galaxies, a new component in the velocity range 7500−8000km s −1 is clearly present in all the four spectra. This emission is not associated with any of the Hα emitting regions since no one of the star forming objects has a recessional velocity below ∼ 8000km s−1 (see Table 10.1). This fact strongly emerges in Fig.10.11 where we compared this spectrum with the mean spectrum obtained from the four pointings. A great fraction of the NW Hα trail described in the previous section lies exactly in the region observed by the BIG-NW pointing, suggesting that the HI emission is probably associated with this feature. This could mean not only that there is neutral hydrogen associated with these structures, but also that their recessional velocity is significantly lower than the mean group velocity (assuming that the low velocity component is associated the Hα trails also in the other three pointings), strongly supporting a ram pressure stripping scenario. Infalling at 1700 km s−1 through the ICM, whose density is ρ ∼ 6 × 10−4 atoms cm−3 at their present peripheral location (A. Finoguenov, private comm.), galaxies in this group will experience 146 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Figure 10.10: The HI spectra obtained for each pointing. 10.3. Results 147 Figure 10.11: Comparison between the combined HI spectrum obtained from the four different Arecibo pointings, and the single pointing on the NW trail. It appears clearly the presence of a low velocity component not associated to the bright galaxies in BIG. 148 10. Galaxy Pre-processing: the blue group infalling in Abell1367 a ram pressure: P = ρv 2 ∼ 3 10−11 [dyn cm−2 ] (10.5) to an order of magnitude higher. Assuming a stellar surface densities σS ∼ 3×10−2 g cm−2 and an interstellar gas surface densities σg ∼ 10−3 g cm−2 , the restoring gravitational force of galaxies is: F = 2πGσS σg ∼ 1.3 10−11 [dyn cm−2 ] Thus the restoring gravitational forces (pressure) at their interiors, are significantly smaller than the ram pressure. In the long run, the increasing ram pressure will fully strip their gaseous material leading to the complete ablation of their interstellar gas, thus suppressing the star formation because of fuel exhaustion. A stripped blob of typical radius R of 2.5 kpc and mass M = 108 M ; might even experience a deceleration, (P − F )πR2 a= = 1.6 10−8 [cm s−2 ] M with a consequent measurable velocity decrease of ∆V = 500 km s−1 in a time as short as 108 yrs, as observed in this case. 10.3.4 The fate of the stripped gas Different predictions are made in the literature for what happens to the gas once it has been stripped. The large extent of the Hα trails and its associated HI gas indicates that it can survive for some 108 yr or even 1 Gyr. This may suggest that evaporation by the ICM is slow, e.g. because the heat flow is saturated and/or that a tangled magnetic field slows down the heat flow into the trail (Vollmer et al. 2001), as observed in the extended HI plume recently discovered in Virgo by Oosterloo & van Gorkom (2005). In spiral galaxies, if the HI column density is above a few times 1020 cm−2 , star formation almost invariably occurs (Boissier et al. 2003). The mean column densities in the trail is ∼ this value, suggesting that it could locally exceed this threshold. Hence, star formation could occur locally in the trails, provided the processes that regulate star formation for a cloud in the ICM are similar to those for gas clouds in spiral galaxies. The Hα emission in the trails could be signature of star formation in act, representing the most extended example of extragalactic star formation ever observed. However we have no evidence of stellar emission from the trails and the dynamical picture of BIG is consistent with the idea that at least part of the gas that has been stripped is just ionized by ram pressure. In this case the plasma density derived in section 10.3.2 implies an exceedingly short recombination time in the ionized trails τr = 1/Ne αa ∼ 2-7 Myr, where αa = 4.2 × 10−13 cm3 s−1 (Osterbrock 1989). Can their exceedingly short recombination time of few Myr be 10.3. Results 149 reconciled with an age between some 108 yr and 1.5 Gyr? We need a mechanism to sustain the ionization along the tail and the presence of the cluster IGM comes to help. The clouds stripped from a galaxy infalling onto the IGM might be kept ionized by X-ray bremsstrahlung emission of the IGM. Following Vollmer et al. (2001) and Maloney et al. (1996) the X-ray ionizing photon flux (φi ) is: φi = ln(0.1/0.0136) FX = 8.3 × 108 FX photons cm−2 s−1 −9 1.6 × 10 1.5 where FX is the X-ray flux. Assuming a total cluster X-ray luminosity of 4 × 1043 erg cm−2 (Donnelly et al. 1998) the X-ray flux at a projected distance of ∼ 125 kpc from the X-ray center (where BIG is observed) is ∼ 2.5 × 10−6 erg cm−2 s−1 and φi results ∼ 2.1 × 103 photons cm−2 s−1 . In equilibrium this gives rise to an ionized column density Ne = φi /αa ne . Using ne = 10−2 cm−3 we obtain Ne ∼ 5 × 1020 cm−2 , consistent with value measured in the ionized tails. This simple calculation shows that the stripped gas can survive in the hostile IGM, being kept ionized by the X-ray photons. 10.3.5 The metal content In order to determine the metal content of the observed emission line knots we followed the same procedure described in Chapter 7. The metallicities obtained from the different methods are shown in Tab.10.4. The uncertainty in the abundances is up to ±0.2dex. All the star-forming regions in BIG are surprisingly metal-rich. Their metallicity lies in the range 8.5 < 12 + log(O/H) < 8.9. It is well known that irregular and spiral galaxies follow a ”metallicity - luminosity relation” (Skillman et al. 1989). Fig.10.12 shows the ”metallicity - luminosity relation” for galaxies in the Virgo cluster (empty circles, taken from Gavazzi et al. in prep., and obtained using the same methods and calibrations) and for the star-forming systems in BIG (triangles). The two bright galaxies CGCG97-114 and CGCG97-125 have a normal metal content for their luminosity. Conversely the star-forming knots show higher abundances for their intrinsic luminosities. If the faintest systems (K1, K5, K2, 114a and 114b) were isolated independently-evolved dwarf galaxies we would measure a metallicity 0.6-1.2 dex lower than the one observed in this case. Moreover their abundances are consistent with the values measured for tidal dwarf systems showing a metal content of 12 + log(O/H) ∼ 8.60 independent from their absolute magnitude (e.g. Duc & Mirabel 1999; Duc et al. 2000). The HII regions DW1, DW2 and DW3 have a high-metal content but consistent, within the calibration uncertainties, with the abundances observed in dwarf galaxies of the same luminosity. However (as discussed in Section10.3.1) the star-forming knots in DW3 are probably not gravitationally 150 Object K1 K1 DW3 a DW3 d 97-114b 97-114a 97-114 K2 a K2 b DW2 a K5 DW1 b DW1 c DW1 c DW1 a DW1 a 97-125 97-125 97-125b 97-125b K3 a: b: c: d: e: 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Tel. Ra23 Rb23 ESO T NG T NG ESO ESO ESO LOI ESO T NG T NG ESO MMT MMT T NG ESO MMT OHP T NG ESO T NG T NG 8.53 8.55 − 8.91 8.59 8.86 9.00 8.56 − 8.53 − 8.56 − 8.81 − 8.55 − − − 8.76 8.89 8.51 8.52 8.62 8.79 8.52 8.71 8.83 8.51 8.64 8.44 − 8.53 8.42 8.70 − 8.52 8.73 8.49 − 8.64 8.75 NII/OIIc NII/Hαd OIII/NIIe Mean Stdev − 8.75 − 8.92 − 8.68 8.85 8.65 − 8.61 − − − 8.84 − 8.67 − 8.77 − 8.75 8.90 − 8.66 8.66 8.73 8.30 8.54 8.70 8.59 8.39 8.67 8.35 8.43 8.64 8.71 8.43 8.58 8.82 8.98 8.45 8.67 8.84 − 8.55 8.70 8.68 8.48 8.72 8.92 8.56 8.64 8.71 8.66 8.48 8.58 8.66 8.49 8.53 8.65 8.72 8.59 8.67 8.78 8.52 8.61 8.66 8.81 8.47 8.70 8.86 8.57 8.56 8.59 8.50 8.50 8.55 8.75 8.46 8.57 8.73 8.74 8.52 8.69 8.83 Zaritsky et al. 1994 McGaugh 1991 Kewley & Dopita 2002 Van Zee et al. 1998 Dutil & Roy 1999 Table 10.4: Metallicities of the galaxies in the BIG group. 0.02 0.09 0.04 0.11 0.13 0.10 0.11 0.05 0.14 0.11 0.22 0.06 0.12 0.08 0.04 0.06 0.08 0.20 0.10 0.05 0.07 10.3. Results 151 Figure 10.12: The relation between Metallicity and B-band Luminosity (with linear best-fit) for galaxies in nearby clusters (empty circles, adapted from Gavazzi et al. 2004). The triangles mark the mean metallicity obtained for the individual knots of BIG. bound, thus each knot should be considered as a single faint extragalactic HII region with a metallicity ∼ 0.8 dex higher than the one obtained from the metallicityluminosity relation. These results rule out an evolutionary scenario in which the faint HII region discovered in BIG are normal independently evolved dwarf galaxies, reinforcing the scenario of Sakai et al. (2002) who proposed that these systems formed from enriched material stripped by tidal interactions from the two brightest galaxies in BIG. 10.3.6 Dating the starburst. Contrary to the gaseous filaments, current star formation is clearly observed in all compact HII regions, dwarf and giant galaxies composing BIG, suggesting that bursts of star formation are presently taking place in this group. Do we have any hint on when the inset of the star bursting phase took place? The dwarf galaxy DW2, and in particular knots DW2b and DW2c show clear Post-Star-Burst signatures in their spectra, with low residual current star-formation. They have an extremely blue continuum (B-R ∼ 0.16), strong Balmer absorption (EW(Hδ)∼ 8Å) and [OII] and Hα in emission. In particular a clear age gradient is observable passing from DW2a, where star formation is still in place, to DW2c, that shows strong Balmer lines in absorption with some evidences of residual star formation (see Fig.10.18). These 152 10. Galaxy Pre-processing: the blue group infalling in Abell1367 features indicates that the starburst ended already ∼ 108 years ago (e.g. Poggianti & Barbaro 1997; Poggianti et al. 1999; Kauffmann et al. 2003b). The star formation history of CGCG97125 The best piece of information for dating the interaction is provided by the brightest group galaxy: CGCG97-125. This galaxy is classified as S0a in the CGCG catalogue (Zwicky et al. 1961), consistent with its red B-R color index (∼ 1.34, see also Fig.10.17) and with the shape of the continuum optical spectrum (Fig.10.13). However this system is far from being a normal early type galaxy. The presence of stellar shells around CGCG97-125 (see Fig.10.6) clearly indicates a past interaction/merger event, as also supported by the disturbed rotation curve analyzed in the previous sections. Numerical simulations predict that the stars from a satellite make a system of shells several 108 yr after the end of the merging event and then they last for more than 1 Gyr (Kojima & Noguchi 1997). The spectrum of 97-125 shows a continuum and absorption features typical of elliptical galaxies; however superimposed to it there are strong emission lines (see Fig.10.13) indicating that this galaxy is still experiencing a strong burst of star formation: a kind of rejuvenated early type galaxy. Using the blue line-strength indices to determine the age of the last star forming event (Longhetti et al. 1999) (Hδ/FeI ∼1.00 , H+K(CaII)∼ 0.91 and ∆4000∼ 1.78) we estimate that the age of the last starburst is ∼1 Gyr, in agreement with the prediction derived from the presence of the stellar shells. However these models assume an instantaneous burst (SSP), that is clearly not the case of CGCG97125, the obtained age thus represents only a lower limit of the real burst age. In Fig.10.13 we compare the drift-scan integrated spectrum and the nuclear spectrum of CGCG97125 obtained at the OHP and TNG telescope respectively: the integrated spectrum appears considerably bluer than the nuclear one. We can use this difference in order to try to reconstruct the recent star formation history of this galaxy. Therefore, assuming that 1)the continuum of the nuclear spectrum is dominated by the old stellar population with no significant contribution from the recent starburst while 2) the integrated one is strongly contaminated by new stars produced during the burst, we can try to estimate the age of the interaction and the stellar mass produced during the burst. Tidal interactions and merging usually produce a sinking of the gas to the galaxy center triggering a burst of star formation, in contrast with our first assumption. Thus in order to test the validity of our hypothesis we used the SED fitting procedure proposed by Gavazzi et al. (2002a) and developed by Franzetti (2005). We assume a ”a la Sandage” star formation history (SFH): SF H(t, τ ) = t2 t × exp(− ) τ2 2τ 2 (10.6) and the Bruzual & Charlot (2003) (BC03) population synthesis models. We fitted 10.3. Results 153 Figure 10.13: Comparison between the drift-scan integrated (blue) and nuclear (red) spectrum of CGCG97-125. Spectrum Nuclear Starburst Z M ass τ t Z log(M/M ) Gyr Gyr 0.04 0.04 11.01 9.27 1.00 0.80 13 1.4 Table 10.5: Best-fitting parameters for the nuclear and starburst component of CGCG97125. 154 10. Galaxy Pre-processing: the blue group infalling in Abell1367 the nuclear spectrum of CGCG97125, corrected for extinction2 assuming t=13 Gyr, a Salpeter IMF (α = 2.35 from 0.1 to 100 M ; Salpeter 1955) and exploring a parameter grid in metallicity (Z) and τ . Z is let free to vary from 1/50 to 2.5 Z in five steps: 0.0004, 0.004, 0.008, 0.02, and 0.05. τ varies from 0.1 to 25 Gyr in 45 approximately logarithmic steps. The best-fitting parameters obtained using the BC03 models are summarized in Table 10.5. The best value of τ is consistent with the one (τ ≤ 3.1Gyr) obtained by Gavazzi et al. (2002a) fitting a template of S0 galaxies. This result validates our assumption that the continuum of the nuclear spectrum is dominated by an old stellar population of the same age expected for an unperturbed S0. By normalizing the obtained model to the observed H-band magnitude and subtracting it to the integrated UV to near-IR SED of CGCG97125, we have the possibility to estimate the starburst contribution to the galaxy emission, the burst age and the stellar mass produced during the star formation. The best-fitting parameters obtained for the starburst SED are summarized in Table 10.5. The burst age results ∼1.4 Gyr and the stellar mass produced during the burst is ∼ 2 × 109 M , consistent with the values previously obtained from independent estimates (i.e. dynamical models). The resulting best fitting SED for CGCG97-125 is shown in Fig.10.14 (black model). The model well reproduces the observations from the far-ultraviolet to the near infrared, with the exception of the near-ultraviolet. This disagreement does not depends on the model assumption but on the attenuation law used to correct for internal dust attenuation. In fact, as shown in Appendix II, we assume a Milky Way attenuation law (thus with a bump at ∼ 2175 Å) that seems not to be valid for normal star forming galaxies (see Chap. 7), producing an overestimate of the real galaxy emission in near-ultraviolet. However this does not influence our results as shown in Fig.10.14. We can thus conclude that the star burst in CGCG97-125 initiated ∼1-1.5 Gyr ago, probably produced by a minor merging of a ∼ 2 × 109 M satellite, and is still taking place. Our result points out that a minor merging able to disturb the morphology and the dynamics of a giant galaxy as CGCG97-125, seems not able to strongly modify the mean age of the stellar population, producing only a small fraction (∼2%) of new stars (see also Fig.10.15). As in the case of NGC4438 (see Chapter 8), this result might probably be representative only for a minor merging into, today gas poor, early type galaxies. 2 We corrected all the spectrophotometric data using the ultraviolet spectral slope β as suggested in Chapter 7 for FUV data and the method described in Appendix A for NUV and optical observations. We assume that the nuclear and integrated spectrum are affected by the same amount of dust extinction, as supported by the similar value for the Hα/Hβ ratio obtained in the two spectra (see Table 10.2). 10.3. Results 155 Figure 10.14: The SED of CGCG97-125, corrected for internal extinction. Nuclear and drift-scan integrated spectra are shown in green. Black circles indicate photometric observations and their relative uncertainties. Best fitting models for the nuclear spectrum (red) and for the starburst component (blue) are given. The resulting best fitting SED for CGCG97-125 is presented in black. Figure 10.15: The star formation history of CGCG97-125 as obtained from the SED fitting procedure. 156 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Figure 10.16: The 2D high resolution spectrum (left) and the optical rotation curve (right) of CGCG97-120 10.3.7 CGCG97-120: simply a foreground galaxy, or an high velocity intruder? In the previous sections I never mentioned the brightest galaxy present in the BIG region: the spiral galaxy CGCG97-120. This massive system has a recessional velocity of ∼5635 km s−1 , thus blueshifted with respect to A1367 by approximately 800 km s−1 . Observations of the neutral hydrogen line show that CGCG97-120 has lost approximately 90% of its original hydrogen content (HI deficiency = 0.9), suggesting that the galaxy has crossed the cluster core and that the ram pressure exerted by the dense intergalactic medium might have caused its hydrogen deficiency. The great velocity difference between this galaxy and BIG (∼ 2500km s−1 ) seems to rule out any association between the two systems, as argued by Sakai et al. (2002) and Gavazzi et al. (2003b). However the deep Hα images obtained at the TNG telescope repropose the question: one of the Hα trail traces in fact a perfect loop around CGCG97-120. Only a blind chance? As shown in Fig.10.6 and Fig.10.16 the galaxy morphology and kinematic are completely unperturbed, showing no signs of interaction. Moreover the scattering angle of an interaction between a satellite galaxy and CGCG97-120 would be of only ∼4-10 degrees, assuming a classical scattering model and an impact parameter of ∼10 kpc: too small to produce the observed loop. Thus for the moment we have to suppose that the association between CGCG97-120 and the Hα trails is only a blind chance. 10.4. Discussion 157 10.4 Discussion 10.4.1 The evolutionary history of the Blue Infalling Group The amount of information provided by the multiwavelength observations presented in this paper allow us to reconstruct the evolutionary history of BIG, during the last 1-2 Gyr. At the beginning of the story BIG was a normal compact group of galaxies with a typical dispersion velocity of ∼150-200 km s−1 , composed of at least three galaxies: a massive evolved early type spiral (CGCG97-125), a massive late type spiral (CGCG97-114) and a gas rich dwarf galaxy (the satellite that has feed CGCG97-125) with a stellar mass ∼ 109 M . Lying in the outskirts of Abell1367 it has been attracted by the cluster potential starting its infall into the cluster core at a mean velocity of ∼ 1700 km s−1 . During their journey, all galaxies are perturbed by gravitational interaction with members, as observed in all compact groups. Stars and gas are stripped, forming tidal tails, bridges (as K2), extragalactic compact HII regions (as K5 and K1) and tidal dwarfs (as DW1, DW2 and DW3). Tidal interactions lowered the restoring force by loosening the potential well of all galaxies in the group making easier stripping gas from the infalling galaxies by ram pressure and producing the unique Hα trails observed in BIG. In particular the gas rich satellite is partially dismantled by the combined action of tidal forces and ram pressure, and finally merged into CGCG97-125 producing stellar shells and a burst of star formation. The combination of gravitational forces and ram pressure is not only consistent with the evidence that BIG is a compact group that is infalling at ∼ 1700 km s−1 into the core of Abell1367, but is also necessary to try to explain all the aspects that make BIG so unique among other known interacting systems and merger remnants: i.e. the unexpectedly high star formation observed in this group, the presence of extended Hα trails and its associated neutral hydrogen, the lack of large-scale tidal tails and, as pointed out by Gavazzi et al. (2003b), the colors of dwarf objects DW1, DW2 and DW3 that are significantly bluer than tidal dwarfs observed in interacting systems (Weilbacher et al. 2000). The IGM compression is in fact able to trigger some star formation in the gas clouds contained within tidal structures (Bekki & Couch 2003), while ram pressure may push some of these clouds free of their parent galaxies, explaining the absence of tidal features and the extremely blue colors of the dwarf objects in BIG. Recently Mayer et al. (2005) have shown that gravitational tides can aid ram pressure stripping by diminishing the overall galaxy potential. The gas stripped along tails fragments into dense clouds and sheet-like structures pressure confined by the ambient medium with the approximately the same column density observed in our case. However their simulations are focused on the evolution of dwarfs (Vrot ∼ 40 km s−1 ) systems orbiting around a Milk Way like galaxy, and it is not clear what would be the effects of the same mechanisms on a massive galaxy infalling into a cluster. 158 10.4.2 10. Galaxy Pre-processing: the blue group infalling in Abell1367 The contribution of preprocessing to cluster galaxies evolution. Galaxy clusters formed not by accreting individual galaxies randomly from the field, but rather through the infall of small groups, falling in along large scale filaments; thus this group represents an unique laboratory reproducing the physical condition expected in a cluster still in formation. What can we learn about galaxy cluster evolution studying BIG? First of all, we are witnessing the first clear example of a well formed S0 galaxy infalling into the core of a cluster of galaxies. This observational evidence suggests that S0 galaxies can form outside clusters and subsequently fall into them: groups environment is in fact considered as the best place where gravitational interactions should operate efficiently and transform a normal spiral into an S0. Moreover gravitational interactions among group’s members are still in act, and CGCG97-125 has recently (∼1.5 Gyr) experienced a minor merging event. The burst of star formation, however, is not able to strongly affect its global optical properties, since the mass of new stars produced is only ∼ 2% of the whole galaxy mass, consistent with the recent results obtained by Boselli et al. (2005a) in the Virgo cluster. This suggests that the mechanism responsible of the transformation of CGCG97-125 into an S0 is older than 2 Gyr, corresponding to a redshift z ≥ 0.2. At its current SFR (1.49 M yr−1 ) the total HI mass of CGCG 97-125 would be depleted in 2.6×109 yr, implying a total burst duration of ∼4 Gyr: consistent with the typical time-scale of the Butcher-Oemler effect (Butcher & Oemler 1978, 1984). Tidal interactions within group members are not only able to produce morphological transformation in galaxies, but also to create new systems formed by gas and stars stripped from group’s members. This is the case of the extremely high number of metal rich star forming dwarfs/extragalactic HII region detected in the infalling group. What will be the future of these stripped systems? It is improbable that all the stripped clusters will infall into the main galaxies, rebuilding the gaseous disk as observed in field mergers: in fact cluster tides and ram pressure stripping act mutually to strip off the material and to inhibit the disk resettling process. If they are dynamically bounded, they could be the progenitors of dwarf cluster galaxies as described in models by Kroupa (1998) and Duc et al. (2004). Simulations predict that young compact massive star clusters formed during the merger of gaseous disk galaxies coalesce within a few 100 Myr forming objects with masses of order 107 − 109 M , as observed in this group, with negligible dark-matter content. However, till today major mergers were supposed not to have produced a significant fraction of the dwarf population, since each merger are expected to spawn only one or two tidal dwarf galaxies. Thus the discovery of a great number of extragalactic star forming knots in BIG (∼10 for one merging event) seems indicate that a tidal formation scenario for part of the dwarf population in cluster could be reasonable, especially at higher redshift, where groups like BIG are expected to infall at higher rate into the core of 10.4. Discussion 159 young clusters. Being produced by a single starburst, these systems might also evolve into dwarf ellipticals, typical of rich clusters. Otherwise, if they will disperse they stars and gas into the cluster their will simply increase the fraction of unbound stars, contributing to the Abell1367 intracluster light, supporting the idea that preprocessing could have had a strong contribution in the amount and distribution of intracluster light (Mihos 2004b). A strong contribution to the intracluster light in Abell1367 would also be provided by the Hα trails if some residual star formation is taking place. In this case these features would represent the most extended example of extragalactic star formation ever observed in the Universe. Surely the Hα trails are strongly contributing to the ICM enrichment, suggesting that a considerable amount of the cluster enrichment might derive from these late-type intruders, as opposed to winds from elliptical galaxies, commonly accepted as the major sources of pollution (e.g. Madau et al. 2001; Mori et al. 2002). This idea is strongly supported by the presence in the NW part of Abell1367 of other two galaxies with Hα trails (Gavazzi et al. 2001a), pointing out that this may not be a rare phenomenon in young clusters. Moreover in the last years an increasing number of X-ray (Hayakawa et al. 2004) and optical (Gavazzi et al. 2001a; Oosterloo & van Gorkom 2005) observations has shown that ram pressure stripping could have an important role on IGM enrichment and recent combined N-body and hydrodynamic simulations have pointed out that more of the 10% of the intracluster medium originated from gas stripped by ram pressure (Domainko et al. 2005). Thus, combined X-ray and optical studies of infalling groups should help us to shed more light on the effect of preprocessing not only on the evolution of cluster galaxies, but also on process of IGM enrichment, an issue which remains unsettled (Tornatore et al. 2004). The evolutionary scenario here presented points out the great importance of groups like BIG not only for galaxy evolution but also for the evolution of clusters itself. Since the infall rate of these groups is considerably higher at high redshift, this analysis points out the strong contribution that small compact groups have probably had in shaping the properties of both galaxies and clusters of galaxies. Thus BIG represents a Rosetta Stone group, giving us the chance to shed light, with the great details possible only in the local Universe, on physical processes typically expected in young, far away, clusters. 160 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Figure 10.17: B-R color map of BIG (Blue = B; Red = R). 10.4. Discussion 161 Figure 10.18: The observed smoothed (step 3) one dimensional spectra. The object identification and telescope are labeled on each panel. 162 10. Galaxy Pre-processing: the blue group infalling in Abell1367 Figure 10.18: Continue. 10.4. Discussion 163 Figure 10.18: Continue. Chapter 11 Discussion & Conclusions 11.1 Discussion In the Introduction to this thesis I have argued that try to recover galaxy evolution during the last 13 Gyr only from observations of today’s Universe represents a real, but fundamental, challenge. The Universe we inhabit is old, and most of the fun is over. In addition (and this is the worst part of the story) the Universe dramatically evolved itself, continuously altering the physical conditions of the environments populated by galaxies. However, as shown in this work, we can still achieve important pieces of information from the study of local galaxies and, combining this information with that obtained at higher redshift, we can try to paint a picture of our knowledge about the evolution of galaxies in clusters; exactly the effort that I’m going to attemp in this conclusion. What evolutionary scenario emerges from this work? The first estimates of the UV cluster luminosity functions from FOCA, FAUST and GALEX observations, here presented, point out that at these wavelengths the cluster LF is considerably steeper than the field one. The steepening of the UV LF from low to high density environment is due to the increasing contribution of early-type, non star forming galaxies, passing from the field to the cluster core. This represents the first evidence of a morphology/star formation - density relation at ultraviolet wavelengths and demonstrates that we cannot blindly consider UV selected galaxies as star-forming systems, especially at low UV luminosities and in high density environments. However this also point out the strong potential of ultraviolet observation in studying all cluster galaxies: not only star-forming systems which UV emission traces the presence of newly born stars, but also early type galaxies in which such emission must be ascribed to low mass old post asymptotic giant branch stars. So let me summarize what I have learned about the evolution of these different morphological types. 165 166 11. Discussion & Conclusions The evolution of elliptical galaxies For the first time, in this work the UV properties of early-type galaxies have been studied down to MB ≤ -15 mag. The newest result addresses the question raised by O’Connell (1999) concerning the dependence of the UV properties on galaxy morphology. We have shown that a dichotomy exists between giant and dwarf ellipticals and, to a lesser extent, between ellipticals and lenticulars. The blueing of the UV color index with luminosity, metallicity, and velocity dispersion indicates that the UV upturn is more important in massive, metalrich systems. Since the UV upturn originates from a minority population of old hot helium-burning horizontal-branch (HB) stars, which emission becoms detectable after at least 10 Gyr (e.g. O’Connell 1999; Brown et al. 2000; Greggio & Renzini 1990; Tantalo et al. 1996), the relation found for giant ellipticals and its small dispersion suggest that clusters ellipticals represent an old, homogeneous population. This is also consistent with the dynamical analysis of Abell 1367 where we found evidence that elliptical galaxies have a Gaussian velocity distribution with a smaller velocity dispersions than the whole cluster sample, representing the virialized, old, cluster population. This picture is supported by both higher redshift observations and N-body simulations. The population of elliptical galaxies in clusters show little evolution in their colors and no structural evolution since at least redshift of z ∼1 (Treu et al. 2005; Smith et al. 2005). The attempt to reproduce this observational evidence with N-body simulations (Springel & Hernquist 2003) results in the invalidation of the paradigm of elliptical formation by mergers of spiral galaxies. At the time the cluster ellipticals were formed in rich clusters, there were simply few if any spiral to merge. It appears clear from simulations, as well as observations of the high-z Universe, that large spiral galaxies as we know them today were very rare at z>2 (Driver et al. 1998; Dickinson et al. 2003; Trujillo et al. 2004; Conselice et al. 2005). It is specially true in dense environment where galaxies had too little time to form large disk from the accretion of high-angular momentum material. This is also supported by HST Deep fields: beyond z=1 the number density of large, well-formed spirals begins to rapidly diminish in favor of a smaller, chaotically arranged systems (Labbé et al. 2003). In addition the recent Millennium simulation (Springel et al. 2005) has shown that the giant ellipticals observed in nearby clusters today, were already formed and massive at very high redshifts (z ∼ 16) and harboured in the center of regions where the first structures developed: the progenitors of rich clusters. We can thus conclude that both observations (at all redshifts) and simulations are consistent with the idea that clusters giant ellipticals are an old, homogeneous population showing no or little evolution at least in the past 8 Gyr. This is not the case for the cluster population of dwarf elliptical galaxies. The opposite behavior in the UV color magnitude relations (reddening of the UV color index with luminosity) of dwarfs with respect to giant ellipticals, similar to that observed for spirals, indicates that the UV spectra of low luminosity objects are shaped by the contribution of young stars, thus presenting a very different star formation history. 11.1. Discussion 167 This implies that the stellar population of dwarfs has been formed in discrete and relatively recent episodes, as observed in other nearby objects (Grebel 2000). However this result is not sufficient to discriminate between different theoretical models for dE formation: primordial objects that lost their gas in a supernova-driven galactic wind (Yoshii & Arimoto 1987; Nagashima & Yoshii 2004), dwarfs irregular infalling into cluster and transformed by ram pressure (van Zee et al. 2004) and/or harassment (Moore et al. 1998), or tidal dwarfs (Kroupa 1998; Duc et al. 2004). The higher frequency of dwarf ellipticals in high density environments supports the idea that they are objects transformed by the harsh cluster environments. However the presence of observational evidence supporting at least the first two scenarios, and the very high dispersions observed in the UV color magnitudes and in structural and kinematic parameters (de Rijcke et al. 2005) seem to suggest that dwarf ellipticals are most likely a mixed population with primordial, and more recently transformed objects co-existing in the present day Universe. Moreover the reddening of the UV color index with luminosity is new evidence that mass drives the star formation history in hot systems (Trager et al. 2000; Gavazzi et al. 2002a; Caldwell et al. 2003; Poggianti 2004b) as in rotating ones (Gavazzi et al. 1996; Boselli et al. 2001, see also below). This phenomenon, today refereed as downsizing effect, is observed in both cluster and field and at least till z ∼0.8 indicating the presence of an ”anti-hierarchical” history for star formation in galaxies. The presence of a downsizing effects in all galaxies, independent from their morphological type, represents today the major challenge for CDM models. The evolution of lenticular galaxies Unlike the rather passive evolution observed in cluster ellipticals, much stronger evolution seems present in the population of cluster S0s. The dispersion observed in the UV color magnitude relation, considerable higher than ellipticals’, bears witness to recent, minor episodes of star formation combined with an old stellar population, as determined also from kinematic and spectroscopic observations (Dressler & Sandage 1983; Neistein et al. 1999; Hinz et al. 2003). This result is consistent with recent studies of stellar population in early type galaxies which found significant differences between the ages of the stellar populations of ellipticals and of the S0 galaxies, supporting the scenario of spirals evolving into S0s (Kuntschner & Davies 1998; van Dokkum et al. 1998; Terlevich et al. 1999; Poggianti et al. 2001; Smail et al. 2001). All these results are supported by the fact that the fraction of S0s in rich clusters has increased significantly since a redshift of z ∼ 1 (Smith et al. 2005), with a corresponding decrease of spiral fraction (Dressler et al. 1997). What is the mechanism responsible for the transformation of a gas rich spiral into a lenticular galaxy? The toy model presented in Chapter 9 has shown that ram pressure alone cannot account for all of the S0 population observed in nearby clusters. 168 11. Discussion & Conclusions Galaxy-cluster IGM interaction is able to remove most of the gas reservoir inducing important structural modifications in the disk properties. We expect that cluster spirals have (at least at short wavelengths) more truncated disk profiles, inverting the outer color gradient with respect to similar but unperturbed objects, and then producing anemic spirals, similar to disk dominated S0s. The surface brightness of the disk, however, mildly decreases in Hα and in the UV bands while remaining mostly constant at longer wavelengths even 5 Gyr after the interaction. Excluding the interaction with the ICM, the only mechanism able to produce a structural modification in spiral galaxies are gravitational interaction. Tidal interactions between galaxies affect both stars and gas. Stars respond by forming arms and bars, while the gas flows directly toward the central regions within about 108 yr after the initial collision. The sinking of the gas towards the galaxy center could trigger a burst of star formation and, on longer timescales, a truncation of the stellar disk (Iono et al. 2004), thus altering galaxy morphology. On the other hand we have to exclude harassment since its influence is largely limited to low luminosity galaxies, while in bright spirals its effects are much more limited (Mihos 2004a; Moore et al. 1996). Thus merger-driven S0 formation mechanisms appear not to work inside the cluster potential, since low velocity interactions are extremely rare. On the other hand, these processes should operate efficiently in the group environment, where the encounter velocities are smaller and cluster tides and the hot ICM are not important. The group environment can create S0s and feed them into the accreting cluster. Although the accretion into the cluster core is expected to happen at higher redshift, we have shown in Chapter 10 a clear example of this phenomenon observed in the local Universe. The starbursting group infalling into the core of Abell1367, represents probably the best example of galaxy preprocessing ever observed. The brightest member of this group is an S0 galaxy in strong gravitational interaction with the other group members. It is thus likely that many of these S0s were processed via mergers in the group environment before being incorporated into clusters; especially in the past where the groups’ infall rate was considerably higher than today. Moreover the discovery of other groups of S0 galaxies in strong interaction such as the one in the outskirts of the Ursa Major cluster strongly support this formation scenario (van Gorkom 2004). This formation scenario is supported by the observational evidence that the bulk of S0 population in clusters was formed between z∼0.2 and z∼1, when the rate of infall of small group was the highest experienced by clusters of galaxies (Mihos 2004a). Finally, S0s are a heterogeneous class, from the bulge dominated to the disky S0s, and it should not be surprising that a single mechanism cannot fully account for the range of S0s types (Hinz et al. 2001, 2003): if ram pressure is able to produce disk dominated S0s (objects similar to the anemic galaxies of Van der Berg), tidal interaction (and thus preprocessing) are required to account for bulge dominated S0s. 11.1. Discussion 169 The evolution of the star formation activity in cluster spiral galaxies We can conclude that the bulk of the bulge dominated S0 cluster population was formed at higher redshifts, and in environments where tidal interactions were more probable. However the morphology density relation, and in particular the star formation density relation, as we observe it today is not fully established at high redshift, because we observe how it evolves in clusters, with star forming late type spirals being transformed into anemic galaxies with quenched star formation of the same morphological type. In order to have an idea of this phenomenon, look at Fig.11.1. It shows the distribution of median value of EW(Hα) as a function of the morphological type for Virgo galaxies with normal gas content (HI-def< 0.4; filled circles) and deficient galaxies (HI-def> 0.4; empty circles). The figure emphasize that, within each Hubble class, galaxies with normal HI content have EW(Hα) systematically higher by a factor two then their deficient counterpart. What is the major mechanism (if any) responsible of this reduction is still unknown: observational results are not always consistent each other, and their interpretation results not straightforward at all. Let me start from the results presented in Chapter 3 and 4. I compared the UV luminosity function of nearby clusters and local field showing that the shape of the LF for star-forming galaxies does not change significantly in different environments. The easiest interpretation of this result is that the dwarf to giant star forming galaxies ratio is independent from the environment; the only thing that changes is the absolute fraction of star forming galaxies (i.e. the normalization of the luminosity function). This is a very simple picture but consistent with the recent work of Balogh et al. (2004) who have shown that the distribution of Hα equivalent widths in star forming galaxies does not depend strongly on the local density, while the fraction of star forming galaxies is a steep function of the local density, in all environments. Understanding the origin of these observed trends is one of the most interesting questions to be answered, since it probably include the key to shed light on the environmental influence on today’s galaxy evolution. First of all, these results seem to suggest that the mechanism that affects the star formation when a galaxy enters a dense environment, must work on a short time scale (≤ 107 − 108 yr), and must affect bright and faint galaxies in the same way, in order to preserve the shape of the luminosity function and of the EW(Hα) distribution. Although this excludes strangulation from playing a major role in galaxy evolution, due to its high time scale (≥1 Gyr), it is not clear which mechanism dominates between ram pressure and tidal interactions. A lot of research groups (i.e. Dressler 2004; Balogh et al. 2004; Goto 2005) have proposed tidal interactions as the major mechanism responsible of galaxy transformation. This idea is mainly supported by the fact that the decreasing of EW(Hα) with local density has approximately the same shape in all environments, from cluster to groups (see Fig. 11.2). Since in groups ram pressure stripping is supposed to be absent (even if Fujita 2004 has shown that ram pressure could be also important in groups), the only mechanisms available to quench 170 11. Discussion & Conclusions Figure 11.1: The distribution of the individual HαE.W. measurements in the Virgo cluster along the Hubble sequence (small dots) and of the median EW(Hα) in bins of Hubble type. Error bars are drawn at the 25th and 75th percentile of the distribution.Filled symbols represent HI-def< 0.4 (unperturbed) objects and open symbols HI-def> 0.4 (HI deficient) galaxies. 11.1. Discussion 171 10 1 10 30 30 20 20 10 10 0 0.1 1 10 0 1 0.1 1 10 Figure 11.2: The star formation rate as a function of density, comparing groups of galaxies with clusters. The upper and lower horizontal dashed lines show the 75% percentile and the median of the equivalent widths. The hashed region shows the relation for the complete sample, while the solid line shows the relation for systems with 500 km s−1 < σ < 1000 km s−1 (left) and σ < 500 km s−1 (right). The dependence on local density is identical irrespective of the velocity dispersion of the whole system. Figure taken from Bower & Balogh (2004). star formation are low velocity encounters: we return once again to the idea that the main mechanism responsible of the evolution of spiral galaxies is preprocessing in little groups. It is indisputable that this simple interpretation rules out the galaxy-ICM interaction; however as shown in this work, in local clusters the role of ram pressure seems significant. First of all in Chapter 5, we have shown that star forming galaxies in Abell1367 have an higher velocity dispersion than the quiescent population. This is observed also in other clusters (Sodre et al. 1989; Stein 1997; Biviano et al. 1997; Adami et al. 1998), and reflects the fact that spirals have an higher velocity dispersion than ellipticals, and a velocity distribution hardly Gaussian (Boselli & Gavazzi 2006). By itself it provides evidence for infall of star forming galaxies into clusters. If a consistent fraction of star forming galaxies are still today passing from low to high density environments, their star formation activity will be soon quenched in order to reproduce the observed trends in luminosity function and EW(Hα) distributions. Since in today’s clusters tidal interaction are less probable, this result supports a ram pressure scenario. In addition van Gorkom (2004) has shown that the velocity dispersion of gas rich galaxies is far from Gaussian contrary to the one of HI deficient galaxies, suggesting that gas rich galaxies that enter the cluster center are likely to 172 11. Discussion & Conclusions Figure 11.3: The ratio of the isophotal Hα and r 0 radii as a function of the HI deficiency for galaxies in the Virgo cluster. be serious affected by interaction with the ICM. This is only the first, and if possible less strong evidence, of the role played by galaxy-intracluster medium interaction. In Chapter 9, we have argued that the population of anemic spirals in clusters, with truncated star forming disks, is produced by ram pressure stripping and that the time scale of the interaction is short (∼100 Myr). In addition a growing number of spiral galaxies are found with unusual morphology in HI, Hα and radio continuum, such as the head tail galaxies CGCG97-073 and CGCG97-079 in Abell1367 (see Chapter 5, Gavazzi et al. 1995, 2001a), NGC4522 (Vollmer et al. 2004b), NGC4388 (Yoshida et al. 2004), NGC4569 (see Chapter 9) in Virgo and CGCG160-055 and CGCG160095 (Bravo-Alfaro et al. 2000, 2001) in Coma. These are prime candidates for ongoing ram pressure stripping. In order to try to determine how important are ICM-ISM interactions for galaxy evolution, as a part of the undergraduate thesis of I.Arosio (Arosio 2005), we analyzed the morphological distribution of galaxies in Virgo and Coma showing that the ratio of the Hα to optical radius correlates with the HI deficiency (see Fig.11.3). This result is consistent with the increase of the fraction of galaxies with truncated star forming toward the center of the Virgo cluster, observed by Koopmann & Kenney (2004a) and 11.1. Discussion 173 Figure 11.4: The clustercentric radial distribution of the individual EW(Hα) measurements in the Virgo cluster. High and low (B-band) luminosity galaxies are given with open and filled dots respectively. Median in bins of 0.5 R/RV ir are given. Error bars mark the 25th and 75th percentile of the distribution. with the prediction of the ram pressure model presented in Chapter 9. In addition the strong correlation between EW(Hα) and HI deficiency observed in nearby clusters (Gavazzi et al. 2002c) completes the ram pressure supporting scenario. To summarize, in nearby clusters we observe galaxies that experience ram pressure stripping: the dominant effect on cluster disk galaxies is a reduction of the star formation rate, which goes hand in hand with the HI deficiency, and for most of the galaxies this seems due to ram pressure. How can we conciliate these results with the universal shape of the EW(Hα) vs density relation presented by Balogh et al. (2004)? This apparent contradiction awaits an explanation. What I think emerges from this work is that the truth lies probably in the middle: it is indisputable that galaxy preprocessing (and in particular tidal interaction) has played a significant role, especially in shaping the properties of bright (giant) galaxies at higher redshift; but at the same time it is unquestionable that we observe today normal ”field-like” galaxies, not affected by any preprocessing, infalling alone (or in 174 11. Discussion & Conclusions very loose groups), for the first time into clusters and on which ram pressure’s effects are clearly evident, as shown in the dynamical study of Abell1367. In part, it might thus be correct to affirm that while tidal interactions have dominated in the past and have probably shaped the morphology-density relation for giant galaxies, ram pressure dominates in today clusters and is surely affecting the star formation history of galaxies but with less influence on their morphology. What is still far from being understood is the downsizing effect (i.e. the correlation of the mean age of stellar populations with the mass). In fact this effect is clearly present in clusters (Gavazzi et al. 2002a; Kodama et al. 2004; Poggianti et al. 2004) where, on the contrary, environmental effects (whichever you prefer) are expected to be more efficient in quenching star formation in dwarfs than in giant galaxies, since gas and stars are less bounded to the galaxy. This effect is clearly evident in Fig.11.4 where is shown the dependence of the EW(Hα) on the cluster-centric distance in the Virgo cluster. While the decline in the star formation rate is clear for giant galaxies (passing from EW(Hα)∼35Å at 2 virial radii to EW(Hα)∼6Å in the cluster center), we do not identify any significant trend for dwarf galaxies. This result could be explained if we assume that a significant replenishment of dwarf galaxies is occurring into rich clusters at the present cosmological epoch. An high infall of dwarf systems is also supported by the fact that the velocity dispersion of dwarf star forming galaxies is considerably higher than the one of high luminosity spiral systems (Adami et al. 1998). Thus understanding how and at which rate galaxies infall and have infalled into cluster represents another important key to shade light on the evolution of star formation activity with clusters. If confirmed, a high infall rate for today’s dwarf galaxies will represent a new interesting challenge for hierarchical models of galaxy evolution, in their unfinished attempt of reproduce the Universe we inhabit. 11.2. Conclusions 11.2 175 Conclusions In this thesis I have investigated the environmental effects on galaxy evolution in nearby clusters using a multiwavelength dataset. In particular this analysis has been focused on the properties of three different local clusters: Abell1367, Virgo and Coma. These three clusters are among the best studied in the local Universe and, due to the variety of their environmental conditions (e.g. spiral fraction, X-ray luminosity, evolutionary stage), they represent the most suitable ”laboratory” for comparative studies. By combining for the first time GALEX UV observations with optical, near and far infrared data we derived extensive observational evidence of cluster galaxy evolution. • I determined the first far-UV and near-UV luminosity functions for nearby clusters finding that in clusters the faint end slope is steeper than in the field. This difference is entirely due to the contribution at low UV luminosities of non-star-forming, massive early-type galaxies that are significantly overdense in clusters; while the luminosity function of cluster star-forming galaxies is consistent with the field one. This indicates that, whatever mechanism affects the star formation activity in late-type cluster galaxies, it influences similarly and with a short time scale the giant and the dwarf components. • I investigated the dynamical state of Abell1367 showing that this cluster is still a young cluster forming at the intersection of large scale filaments. At least two subgroups are currently infalling into the main cluster. They show a higher fraction of star forming galaxies than the cluster core, as expected during the early phase of merging events, confirming that the building up of large scale structures can strongly affects the evolutionary history of galaxies. • I studied for the first time the UV properties of a volume-limited sample of early-type galaxies showing the presence of a clear dichotomy in the FUV-optical color magnitude relations between giant and dwarf ellipticals. For elliptical and lenticular galaxies, the (FUV-NUV) color becomes bluer with increasing luminosity and with increasing reddening of the optical or near-IR color indices. For the dwarfs, the opposite trend is observed. These results are consistent with the idea that the UV emission is dominated by hot, evolved stars in giant systems, while in dwarf ellipticals residual star formation activity is more common. • While investigating the star formation history of galaxies in nearby clusters using ultraviolet observations, it has been mandatory to study UV dust attenuation properties of nearby galaxies in order to look for new recipe’s in oder to correct GALEX data. I confirmed that normal galaxies follow a LT IR /LF U V − β relation offset from the one observed for starburst galaxies. The dispersion of 176 11. Discussion & Conclusions this relation is found to weakly correlate with the galaxy star formation history. I studied the correlation of dust attenuation with other global properties, such as the metallicity, dynamical mass, ionized gas attenuation, Hα emission and mass surface density providing some empirical relations from which the total infrared to far ultraviolet ratio (LT IR /LF U V ) can be estimated when far infrared data are absent. This result represents only the tip of the iceberg of a study of dust properties in normal galaxies. Only comparing data with models we will be able to properly correct data for dust extinction and thus to estimate the star formation rate in galaxies. Finally I studied in great details the star formation history of three different systems considered as the prototypes of the three main environmental mechanisms able to perturb galaxy evolution, namely: high velocity interactions, ram pressure stripping and galaxy preprocessing. • We showed that in today’s cluster galaxies high-velocity tidal encounters between two galaxies of similar mass are able to perturb the stellar distribution and thus produce important tidal tails, but are not sufficient to significantly increase the star formation activity of cluster galaxies. • Moreover we demonstrated that ram pressure stripping alone is not able to transform a spiral galaxy into an S0, reproducing the structural properties of present-day lenticulars. • Strong transformations in both morphology and star formation activity can be produced by the mutual effects of low velocity encounters and ram pressure stripping in small groups infalling into the cluster core (preprocessing), as observed in Abell1367. Studying this unique example of preprocessing in the local Universe we showed that infalling groups could have a strong influence not only on galaxy evolution but also on the evolution of cluster galaxies, significantly contributing to the enrichment of the intracluster medium and to the intracluster light. Considering all these observational results I conclude that • Giant ellipticals are an old, homogeneous population showing no or little evolution at least in the past 8 Gyr unlike dwarf ellipticals which still contains young stellar populations. • The importance of different environmental mechanism is directly linked with the age of the Universe. • Tidal interactions and prepocessing have probably dominated in the past Universe and has shaped part of the morphology-density relation during the cluster accretion of small groups. 11.2. Conclusions 177 • Ram pressure dominates in today clusters and is surely affecting the star formation history of galaxies with less influence on their morphology. • The heterogeneous class of S0s galaxies, from bulge dominated to the disky S0s, is not the result of a single transformation mechanism: if ram pressure is able to produce disk dominated S0s, tidal interaction (and thus preprocessing) are required to account for bulge dominated S0s. • Different observational clues confirm the presence of a correlation between the mean age of stellar populations and the mass of their parent galaxies (downsizing effect). In the framework of the hierarchical model of galaxy formation, the origin of the downsizing effect remains unsolved and represents one of the main challenges for models of galaxy evolution. Appendix A The extinction correction Here we present the method used to correct for dust attenuation multiwavelength observations used in this work, and proposed by Boselli et al. (2003a) As discussed in Chapter 7 the observed stellar radiation of galaxies, from UV to near-IR wavelengths, is subject to internal extinction (absorption plus scattering) by the interstellar dust. Estimating the dust extinction at different λ in external galaxies is very difficult (it has been done only for the Magellanic clouds). This difficulty is mainly due to two reasons: a) the extinction strongly depends on the relative geometry of the emitting stars and of the absorbing dust within the disc of galaxies. The young stellar population are mostly located along the disc in a thin layer, while the old populations forms a thicker layer. This point is further complicated by the fact that different dust components (very small grains, big grains etc.), which have different opacities to the UV, visible or near-IR light, have themselves different geometrical distributions both on the large and small scales. b) it is still uncertain whether the Galactic extinction law is universal, or if it changes with metallicity and/or with the UV radiation field. Detailed observations of resolved stars in the Small Magellanic Cloud by Bouchet et al. (1985) indicate that the extinction law in the optical domain is not significantly different from the Galactic one in galaxies with a UV field ∼10 times higher and a metallicity ∼10 times lower than those of the Milky Way. A steeper UV rise and a weaker 2200 Åbump than in the Galactic extinction law have been however observed in the LMC and SMC (Mathis 1990). After the results of Chapter 7 the adoption of the Galactic extinction law for external galaxies could seem not reasonable, however in this moment we have not yet a good alternative, moreover no simple analytic functions describing the geometrical distribution of emitting stars and absorbing dust, both on small and large scales, are yet available. The radiative transfer models of Witt & Gordon (2000) have however shown that the FIR to UV flux ratio, being mostly independent of the geometry, of the star formation history (the two radiations are produced by similar stellar populations) and of the adopted extinction law, is a robust estimator of the dust extinction at UV wavelengths. From the value of the TIR/FUV 179 180 A. The extinction correction (measured or obtained with the recipes presented in Chapter 7 we can thus estimate A(FUV), following Buat et al. (2005): A(F U V ) = −0.0333 ∗ y 3 + 0.3522 ∗ y 2 + 1.1960 ∗ y + 0.4967 [mag] (A.1) where y is log(T IR/F U V ). A(λ) can be derived from A(FUV) once an extinction law and a geometry for the dust and star distribution are assumed. We adopt the sandwitch model, where a thin layer of dust of thickness ζis embedded in a thick layer of stars: 1 − ζ(λ) 1 + e−τ (λ)·sec(i) + A(λ) = −2.5 · log 2 ζ(λ) −τ (λ)·sec(i) + [mag] (A.2) · 1−e τ (λ) · sec(i) where the dust to stars scale height ratio ζ(λ) depends on λ (in units of Å) as: ζ(λ) = 1.0867−5.501 × 10−5 · λ. (A.3) This has been calibrated adopting the average between the optically thin and optically thick cases with λ dependent dust to star scale height ratios given by Boselli & Gavazzi (1994). Observations of some edge-on nearby galaxies show that it is still unclear whether ζ depends or not on λ (Xilouris et al. 1999). As shown in Gavazzi et al. (2002a), however, similar values of Ai (λ)are obtained in the case of a sandwitch model and of the extreme case of a slab model (ζ = 1), meaning that the high uncertainty on ζ is not reflected on A(λ). In the case of the FUV band ( λ ∼ 1530 Å), ζ = 1, and Eq. A.2 reduces to a simple slab model. In this case τ (UV) can be derived by inverting Eq. A.2: τ (UV) = [1/sec(i)] · 0.0259 + 1.2002 × Ai (FUV) + 1.5543 × Ai (FUV)2 + − 0.7409 × Ai (FUV)3 + 0.2246 × Ai (FUV)4 (A.4) using the galactic extinction law k(λ) (Savage & Mathis 1979), we than derive: τ (λ) = τ (UV) · k(λ)/k(UV) and we compute the complete set of Ai (λ) using Eq. A.2. (A.5) Appendix B Estimate of the < 912Å flux from Hα + [NII] The stellar radiation field with λ <912 Åionizes the gas, which re-emits, via recombination lines. If the gas is optically thick in the Lyman continuum, the number of photons in a specific recombination line is directly proportional to the number of star photons in the Lyman continuum. In the case of Hβ this number is given by equation (5.23) in Osterbrock (1989). 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P. 1994, ApJ, 420, 87 Zibetti, S., Gavazzi, G., Scodeggio, M., Franzetti, P., & Boselli, A. 2002, ApJ, 579, 261 Zwicky, F., Herzog, E., & Wild, P. 1961, Catalogue of galaxies and of clusters of galaxies (Pasadena: California Institute of Technology (CIT)) List of Figures 1.1 2.1 2.2 2.3 3.1 3.2 3.3 3.4 4.1 4.2 4.3 An example of the heterogeneous population of galaxies that inhabit our Universe. Mosaic of RGB (g,r,i) images adapted from Frei et al. (1996) . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 6 Cross section of the instrument portion of GALEX. The optical path is outlined in blue. Overall dimensions of the view shown are 1.5 m 1 m (adapted from Morrissey et al. 2005). . . . . . . . . . . . . . . . . 16 The transmittance profile for the NUV and FUV GALEX filters. Different galaxy spectral energy distributions are superposed. . . . . . . 19 Example of GALEX image. GALEX NGS observation of NGC4631. In the color table, red-green (gold) is used for NUV, and blue for FUV. 19 The UV luminosity functions for the four analyzed data sets. . . . . . The composite UV luminosity function of 3 nearby clusters. The solid line represents the best Schechter fit to the data for MUV ≤ −16.5. . . The UV bi-variate composite luminosity functions of nearby clusters. Red (UV − B > 2) and blue (UV − B < 2) galaxies are indicated with empty and filled circles respectively. . . . . . . . . . . . . . . . . . . . The cluster and the field UV luminosity functions. The composite cluster LF is given with filled circles. The solid line indicates the best Schechter fit of the field LF of Sullivan et al. (2000). The normalization is such that the two LFs match at MUV ∼ −19.25. . . . . . . . . . . . 27 28 29 31 The GALEX observation of Abell1367. ROSAT X-ray contour are superposed in black. The tick rectangular region indicates the region covered by the optical catalogues used for the star/galaxy discrimination. 34 Comparison between FOCA (upper image) and GALEX (lower image) observation of the center of Abell1367. It emerges clearly the strong improvement in resolution and sensitiveness of new GALEX data. . . 35 Left: The comparison between FOCA and GALEX NUV (left) and FUV (right) magnitudes of galaxies in Abell1367. The continuum line shows the best linear fit to the data. . . . . . . . . . . . . . . . . . . 36 201 202 4.4 4.5 4.6 4.7 4.8 5.1 5.2 5.3 5.4 5.5 LIST OF FIGURES The redshift completeness per bin of UV magnitude in Abell 1367. . . The GALEX NUV (left) and FUV (right) LF for Abell 1367. Open dots are obtained using the subtraction of field counts obtained by Xu et al. (2005); filled dots are obtained using the completeness corrected method. The solid line represents the best Schechter fit. The dotted line shows the composite nearby clusters 2000 Å LF by (Cortese et al. 2003a). The dashed line represents the GALEX local field LF (Wyder et al. 2005), normalized in order to match the cluster LF at MAB ∼ −17.80. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . The NUV (left) and FUV (right) bi-variate LF of A1367. Star-forming and quiescent galaxies are indicated with empty triangles and filled squares respectively. The dashed line represents the GALEX local field LF (Wyder et al. 2005), normalized as in Fig.4.5 . . . . . . . . . The FUV-NUV color magnitude relation for confirmed members of A1367. Symbols are as in Fig.4.6 . . . . . . . . . . . . . . . . . . . . The optical (r 0 -band) distribution for star forming (blue histogram) and quiescent (red histogram) galaxies in our sample. . . . . . . . . . 37 Cumulative redshift distribution for galaxies in the studied region. . . Velocity histogram and stripe density plot for the members of Abell 1367. Arrows mark the location of the most significant weighted gaps in the velocity distribution. . . . . . . . . . . . . . . . . . . . . . . . Local deviations from the global kinematics for galaxies in Abell 1367 as measured by the Dressler & Shectman (1988) test. Galaxies are marked with open circles whose radius scales with their local deviation δ from the global kinematics. The ROSAT X-ray contours are shown with dotted lines. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . Palomar DSS image of the central region (∼1.3 square degrees) of Abell 1367 studied in this Chapter. The iso-density contours for the 146 confirmed cluster members are superposed. The lowest iso-density contour correspond to 3σ above the mean density in the field (left). The ROSAT X-ray contours are superposed in red (right). The straight line indicates the position of the abrupt gas temperature gradient detected by ASCA (Donnelly et al. 1998), used to divide our sample into two subclusters: the North-West and the South-East. . . . . . . . . . . . The LOS velocity field (left) and the velocity dispersion field (right) for the whole region studied in this Chapter. The LOS velocity and the velocity dispersion are computed using the 10 nearest neighbors to each pixel, whose size is 36 arcsec2 . The iso-density contours for the 146 confirmed cluster members are superposed in black. . . . . . . . 46 38 38 41 41 47 51 52 53 LIST OF FIGURES 5.6 5.7 5.8 5.9 5.10 5.11 5.12 5.13 6.1 6.2 A 3D sketch of Abell 1367 summarizing the various sub-components described in Section 5.5. The cluster is viewed from its near side, as suggested by the eyeball indicating the observer’s position. . . . . . . Blow-up of the NW substructure of Abell 1367. The arrows indicate the direction of radio head tails associated with 97-079 and 97-073 and the orientation of the NAT radio galaxy 97-095. The dashed region shows the distribution of the diffuse cluster radio relic (Gavazzi 1978). The iso-density contours for the 146 confirmed cluster members are superposed. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . The LOS velocity distribution for galaxies in the NW (upper) and in the SE (lower) subclusters. . . . . . . . . . . . . . . . . . . . . . . . . The velocity dispersion radial profile of the NW (upper) and the SE (lower) subclusters. . . . . . . . . . . . . . . . . . . . . . . . . . . . . The distribution of galaxies belonging to the South-East subcluster. Triangles indicate galaxies with LOS velocity > 7500 km s−1 , circles galaxies with LOS velocity < 5800 km s−1 and squares galaxies with LOS velocity comprises in the range 5800 km s−1 < V < 7500 km s−1 . The ROSAT X-ray contours are shown. . . . . . . . . . . . . . . . . . The LOS velocity distribution for emission line (upper) and non emission line galaxies (lower) in the whole cluster sample. . . . . . . . . . Projected density map of non emission line (left) and emission line (right) galaxies in Abell 1367. The iso-density contours of the 146 confirmed cluster members are superposed. . . . . . . . . . . . . . . . The bound and unbound orbit regions in the (Vrel , α) plane. The bound-incoming solutions (BIa and BIb ), the bound-outgoing solutions (BO) and the unbound-outgoing (UO) solutions are indicated with solid lines. The dotted lines show the dividing line between bound and unbound regions. The vertical solid lines represent the observed Vrel and the dashed regions their associated 1σ uncertainty. . . . . . The near-UV (left column) and far-UV (right column) to optical and near-IR color magnitude relations. Colors are in the AB magnitude system. Open circles are for ellipticals, filled circles for dwarfs, crosses for lenticulars (S0-S0a). Galaxies redder than the dashed line are undetectable by the present survey (at the NGS limit). Largest 1σ errors for luminous and dwarf systems are given. . . . . . . . . . . . . . . . The relationship between the UV color index (F U V − N U V ) and a) the total H band luminosity, b) the B-H color index, c) the logarithm of the central velocity dispersion and d) the Mg2 index. Symbols are as in fig. 6.1. Labeled points indicate objects having unusual radio or optical properties (see Sect. 3). . . . . . . . . . . . . . . . . . . . . . 203 54 55 56 57 58 59 60 63 72 74 204 6.3 6.4 6.5 7.1 7.2 7.3 7.4 7.5 LIST OF FIGURES The relationship between the UV color index (F U V − N U V ) and the total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for dwarf ellipticals are presented. . . . . . . . . . . . . . . . 77 The relationship between the UV color index (F U V − N U V ) and the total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for ellipticals are presented. . . . . . . . . . . . . . . . . . . 78 The relationship between the UV color index (F U V − N U V ) and the total H band luminosity. Symbols are as in fig. 6.1. The optical spectra available for lenticulars are presented. . . . . . . . . . . . . . . . . . . 79 Ratio of the total infrared to far ultraviolet luminosity as a function of the ultraviolet spectral slope (lower x-axis) and the FUV-NUV color (upper x-axis). Open circles indicates our secondary sample while filled circles represent the primary sample. The dashed line represents the best linear fit to starburst IRX-UV relation. The solid line indicates the best bisector linear fit for our primary sample. The stars indicate the sample of IUE starbursts. Mean error bars for the plotted data are shown in the lower right corner, in this and subsequent figures. The residuals from the best linear fit for normal galaxies are shown in the bottom panel. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 88 Relation between the birthrate parameter computed from the Hα emission, and the distance from the LT IR /LF U V − β relation for starbursts. The solid line represents the best linear fit. . . . . . . . . . . . . . . . 90 The relation between the ultraviolet spectral slope β and the Hα attenuation obtained from the Balmer decrement. Symbols are as in Fig.7.1. Solid line represents the best linear fit to our primary sample (equation 7.14) while the dashed line indicate the best-fit for starburst galaxies obtained by Calzetti et al. (1994) (equation 7.13). Arrows indicate galaxies for which the value of A(Hα) is a lower limit of the real value (i.e. Hβ undetected). The residuals from the best linear fit for normal galaxies are shown in the bottom panel. . . . . . . . . . . . . 93 Relation between gas metallicity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig.7.1. The solid lines show the best linear fit for our primary sample. The residuals from the best linear fits for normal galaxies are shown in the upper panels. . . . . . . . . . . . . . 95 Relation between the galaxy size and the LT IR /LF U V ratio for starburst (left panel) and normal galaxies (right panel). Symbols are as in Fig. 7.1. Mean values and uncertainties in bins of 0.30 log(Diameter) are given. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 95 LIST OF FIGURES 7.6 7.7 7.8 7.9 205 Relation between the gas to dust ratio and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. . . . . . . . . . . . . . . . . . . . . 97 Relation between the H-band luminosity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panel. . . . . . . 98 Relation between the TIR+FUV luminosity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. . . . . . . . . . . . . . 98 Relation between the mean H-band surface brightness (µe ) and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panel.100 7.10 Relation between the star formation rate density and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. The solid line shows the best linear fit for our primary sample.The residuals from the best linear fit for normal galaxies are shown in the upper panel. . 100 7.11 Relation between the Hα and far ultraviolet luminosity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. Hα luminosity is corrected for dust attenuation using the Balmer decrement, while the FUV flux is uncorrected. The solid lines show the best linear fit for our primary sample. The residuals from the best linear fit for normal galaxies are shown in the upper panels. . . . . . . . . . . . . . . . . . 102 7.12 Relation between the observed Hα and far ultraviolet luminosity and the LT IR /LF U V ratio (left) or β (right). Symbols are as in Fig. 7.1. Hα luminosity is the observed value not corrected for dust attenuation. The solid lines show the best linear fit for our primary sample.The residuals from the best linear fit for normal galaxies are shown in the upper panels. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 102 8.1 The combined NUV and FUV image of NGC 4438. The regions described in sect. 3 of the text are labeled 1 to 7. The horizontal line is 10 kpc long (assuming a distance of 17 Mpc). . . . . . . . . . . . . . 109 8.2 The Hα+[NII] contours (red, in arbitrary scale, in between 8 10−17 and 6 10−16 erg cm−2 s−2 arcsec−2 , with σ= 5 10−17 erg cm−2 s−2 arcsec−2 , from Boselli & Gavazzi (2002)) are superposed to the NUV gray-level image of NGC 4438. . . . . . . . . . . . . . . . . . . . . . . . . . . . 111 8.3 Chandra image of NGC 4438 in the 0.3-2 keV energy band with Hα contours superposed. Adapted from Machacek et al. (2004) . . . . . . 111 206 LIST OF FIGURES 8.4 The RGB (FUV=blue, NUV=green, B=red) color map of NGC 4438 and NGC 4435. The SED of each region defined in Fig. 1 are given in the lower plot of each frame. Crosses indicate the observed data, arrows upper limits (in mJy), the red dashed line the evolved population fit as determined by the model of Boissier & Prantzos (2000), the dotted blue line the starburst SED (from Starburst 99) and the dashed green line the combined fitting model. The burst luminosity contribution (for the age corresponding to the minimum χ2 ) in the band FUV, B and K is also given. The upper panel gives the variation of the reduced χ2 parameter (black continuum line, in logarithmic scale) and of the burst mass fraction (red dotted line) as a function of the age of the burst (in Myr). The lower panel of region 4 gives the integrated 3500 to 7000 Å, R=1000 spectrum of the main body of the galaxy (black continuum line) compared to the fitted model (red dashed line). . . . 115 9.1 The radial profile of observed (open symbols) and extinction-corrected (filled symbols) H-band surface brightness (left) and of the rotational velocity (center) used to constrain the model without interaction (represented by the black solid line). The total gas radial profile (right) predicted by the unperturbed model (solid black line) is compared to the observed one (green filled circles), obtained by summing the HI component (red line) to the molecular one (blue and light blue) and correcting for Helium contribution (× 1.4), and to the model including the interaction (black dashed line). . . . . . . . . . . . . . . . . . . . 120 9.2 Ram pressure stripping intensity (in arbitrary units) as a function of time (Eq.9.1). Adapted from Vollmer et al. (2001). . . . . . . . . . . 121 9.3 The radial profile of the observed (empty green circles) and extinctioncorrected (filled green circles) total gas, Hα, FUV (1530 Å), NUV (2310 Å), B and i surface brightness. The yellow shaded area marks the range in between the observed (bottom side) and extinction-corrected (top side) surface brightness profiles. Surface brightnesses are compared to the model predictions without interaction (black solid line) or with interaction for several 0 and t0 parameters. Equal maximum efficiency (0 =1.2 M kpc−2 yr−1 ) and different age: t0 =100 Myr, red continuum line (the adopted model); t0 =500 Myr, grey long dashed line, t0 =1.5 Gyr, dashed magenta line. Equal age (t0 =100 Myr) and different maximum efficiency: 0 =3 M kpc−2 yr−1 , blue dotted line; 0 =1/3 M kpc−2 yr−1 , orange dotted line. . . . . . . . . . . . . . . . . . . . . . 122 LIST OF FIGURES 9.4 9.5 207 The observed and model surface brightness (a), color (b) radial profiles of NGC 4569. In the model profiles the continuum lines are for models with gas removal, dashed lines for unperturbed models. c) the reduced χ2 as a function of the look-back time of the ram-pressure event for a few efficiencies 0 (M kpc−2 yr−1 ). Models were computed each 100 Myr for 0 < t0 < 500 Myr, 200 Myr for 1.5 < t0 < 0.5 Gyr and 1 Gyr for 6.5 < t0 < 1.5 Gyr; and each 0.2 M kpc−2 yr−1 efficiencies between 0.4 and 1.6 (only the more relevant are shown here). d) the variation of the effective surface brightness (mean surface brightness within Re , the radius containing half of the total light) and radius due to differential variation of the star formation history of NGC 4569. Open triangles are for the unperturbed model, the other symbols for different ages of the interaction (100 Myr, 1.5 and 5.5 Gyr). . . . . . . . . . . . . . . . 125 The RGB (continuum subtracted Hα =blue, FUV=green, NUV=red) color map of NGC 4569 . . . . . . . . . . . . . . . . . . . . . . . . . 126 10.1 The four Arecibo HI pointings obtained in the region of the BIG group, superposed to the r 0 band image. The size of each circle correspond to the telescope beam. . . . . . . . . . . . . . . . . . . . . . . . . . . . . 129 10.2 GALEX NUV image of the Blue Infalling group (BIG). . . . . . . . 130 10.3 High-contrast Hα+[NII] band frame of the BIG group. . . . . . . . . 132 10.4 Upper panel: The position and the width (rectangular areas on the right) of the three slits obtained for CGCG97-125. The slits are superposed to the Hα + [NII] net image. Lower Panel: The three different rotations curves obtained for CGCG97-125. Letters indicate the different regions as labeled in the upper panel. . . . . . . . . . . . . . . 136 10.5 The low resolution 2D spectrum obtained at ESO/3.6 for the knots DW3d (left) and DW3e (right), shows a significant difference (∼ 500km s−1 ) in the velocity of the two knots. . . . . . . . . . . . . . . . . . . . . . 137 10.6 Stellar shells are seen around galaxy 97-125 in the r 0 band image of BIG. No continuum emission is detected from the low brightness trails (except K2). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 139 10.7 Extended low brightness trails appear in the Hα+[NII] NET frame of BIG. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 140 10.8 HI column density distribution in BIG. Contours are 0.5, 1.0, 2.0, 3.0, 4.0, 5.0, and 6.020 cm−2 . Adapted from Sakai et al. (2002) . . . . . . 144 10.9 HI position-velocity diagram centered on CGCG 97-125. Adapted from Sakai et al. (2002) . . . . . . . . . . . . . . . . . . . . . . . . . . . . 144 10.10The HI spectra obtained for each pointing. . . . . . . . . . . . . . . 146 208 LIST OF FIGURES 10.11Comparison between the combined HI spectrum obtained from the four different Arecibo pointings, and the single pointing on the NW trail. It appears clearly the presence of a low velocity component not associated to the bright galaxies in BIG. . . . . . . . . . . . . . . . . . . . . . . 10.12The relation between Metallicity and B-band Luminosity (with linear best-fit) for galaxies in nearby clusters (empty circles, adapted from Gavazzi et al. 2004). The triangles mark the mean metallicity obtained for the individual knots of BIG. . . . . . . . . . . . . . . . . . . . . . 10.13Comparison between the drift-scan integrated (blue) and nuclear (red) spectrum of CGCG97-125. . . . . . . . . . . . . . . . . . . . . . . . . 10.14The SED of CGCG97-125, corrected for internal extinction. Nuclear and drift-scan integrated spectra are shown in green. Black circles indicate photometric observations and their relative uncertainties. Best fitting models for the nuclear spectrum (red) and for the starburst component (blue) are given. The resulting best fitting SED for CGCG97125 is presented in black. . . . . . . . . . . . . . . . . . . . . . . . . . 10.15The star formation history of CGCG97-125 as obtained from the SED fitting procedure. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.16The 2D high resolution spectrum (left) and the optical rotation curve (right) of CGCG97-120 . . . . . . . . . . . . . . . . . . . . . . . . . . 10.17B-R color map of BIG (Blue = B; Red = R). . . . . . . . . . . . . . . 10.18The observed smoothed (step 3) one dimensional spectra. The object identification and telescope are labeled on each panel. . . . . . . . . . 10.18Continue. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.18Continue. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 147 151 153 155 155 156 160 161 162 163 11.1 The distribution of the individual HαE.W. measurements in the Virgo cluster along the Hubble sequence (small dots) and of the median EW(Hα) in bins of Hubble type. Error bars are drawn at the 25th and 75th percentile of the distribution.Filled symbols represent HI-def< 0.4 (unperturbed) objects and open symbols HI-def> 0.4 (HI deficient) galaxies. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 170 11.2 The star formation rate as a function of density, comparing groups of galaxies with clusters. The upper and lower horizontal dashed lines show the 75% percentile and the median of the equivalent widths. The hashed region shows the relation for the complete sample, while the solid line shows the relation for systems with 500 km s−1 < σ < 1000 km s−1 (left) and σ < 500 km s−1 (right). The dependence on local density is identical irrespective of the velocity dispersion of the whole system. Figure taken from Bower & Balogh (2004). . . . . . . . 171 LIST OF FIGURES 209 11.3 The ratio of the isophotal Hα and r 0 radii as a function of the HI deficiency for galaxies in the Virgo cluster. . . . . . . . . . . . . . . . 172 11.4 The clustercentric radial distribution of the individual EW(Hα) measurements in the Virgo cluster. High and low (B-band) luminosity galaxies are given with open and filled dots respectively. Median in bins of 0.5 R/RV ir are given. Error bars mark the 25th and 75th percentile of the distribution. . . . . . . . . . . . . . . . . . . . . . . . . 173 List of Tables 2.1 Selected Performance Parameters (Morrissey et al. 2005) . . . . . . . 17 3.1 3.2 Integral redshift completeness in bin of 0.5 magnitudes. . . . . . . . . The completeness-corrected differential number of galaxies per bin of magnitude . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 26 4.1 Best Fitting Parameters. . . . . . . . . . . . . . . . . . . . . . . . . . 37 5.1 5.2 5.3 45 48 5.4 5.5 5.6 5.7 5.7 The spectrograph characteristics . . . . . . . . . . . . . . . . . . . . . 1D substructure indicators for the whole cluster sample . . . . . . . . The most significant weighted gaps detected in the velocity distribution of the whole cluster sample. . . . . . . . . . . . . . . . . . . . . . . . 3D substructure indicators for our sample . . . . . . . . . . . . . . . Mass estimate for Abell 1367 . . . . . . . . . . . . . . . . . . . . . . . Two-body model parameters . . . . . . . . . . . . . . . . . . . . . . . The 119 new redshift measurements . . . . . . . . . . . . . . . . . . . Continue . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 6.1 Main relations for early type galaxies . . . . . . . . . . . . . . . . . . 73 7.1 Linear realtions useful to estimate the LT IR /LF U V ratio (log(LT IR /LF U V ) = a × x + b). . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 105 10.1 Redshifts of the galaxies in the BIG group. . . . . . . . . . . . . . . 10.2 Line fluxes, corrected for internal extinction, of the galaxies in the BIG group. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 10.3 Properties of galaxies in BIG. . . . . . . . . . . . . . . . . . . . . . . 10.4 Metallicities of the galaxies in the BIG group. . . . . . . . . . . . . . 10.5 Best-fitting parameters for the nuclear and starburst component of CGCG97125. . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 211 30 49 50 61 65 67 68 133 135 138 150 153