Download environmental effects on galaxy evolution in nearby clusters

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
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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
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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
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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
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3.4
Local
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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
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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 .
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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 . . . . . . . . . . .
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with GALEX.
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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 . . . . . . . . . . . . . . . . . . .
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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.
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11 Discussion & Conclusions
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11.1 Discussion . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 165
11.2 Conclusions . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . . 175
A The extinction correction
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B Estimate of the < 912Å flux from Hα + [NII]
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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). For Hα we have:
Z∞
Lν
dν = LHα · C
hν
(B.1)
ν0
where:
ef f
αHβ
(H o , T ) FHα
1/C = hνHβ
αB (H o , T ) FHβ
(B.2)
Assuming T =10000K and the Osterbrock case B:
ef f
(H o , T ) = 3 × 10−14 (cm3 sec−1 )
αHβ
αB (H o , T ) = 2.59 × 10−13 (cm3 sec−1 )
= 2.87
and FFHα
Hβ
From the Hα luminosity it is thus possible to recover the number of ionizing photons,
which can be compared with the similar quantity derived from the integral on the
model spectrum.
A conservative estimate of the uncertainty on the derived < 912 Å flux is 1 mag.
181
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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