* Your assessment is very important for improving the workof artificial intelligence, which forms the content of this project
Download AST4930 Star and Planet Formation
Aries (constellation) wikipedia , lookup
Canis Minor wikipedia , lookup
Corona Borealis wikipedia , lookup
Space Interferometry Mission wikipedia , lookup
Theoretical astronomy wikipedia , lookup
Constellation wikipedia , lookup
Nebular hypothesis wikipedia , lookup
Corona Australis wikipedia , lookup
Auriga (constellation) wikipedia , lookup
International Ultraviolet Explorer wikipedia , lookup
Cassiopeia (constellation) wikipedia , lookup
Timeline of astronomy wikipedia , lookup
Observational astronomy wikipedia , lookup
Planetary system wikipedia , lookup
Cygnus (constellation) wikipedia , lookup
Globular cluster wikipedia , lookup
Aquarius (constellation) wikipedia , lookup
Star catalogue wikipedia , lookup
Orion (constellation) wikipedia , lookup
Perseus (constellation) wikipedia , lookup
Corvus (constellation) wikipedia , lookup
Cosmic distance ladder wikipedia , lookup
Future of an expanding universe wikipedia , lookup
Stellar classification wikipedia , lookup
Stellar evolution wikipedia , lookup
Open cluster wikipedia , lookup
Orion Nebula wikipedia , lookup
Stellar clusters and associations • Stars do not generally form in isolation. • The vast majority (90%) of stars form in aggregates with 100 or more members, for a total stellar mass exceeding 50M NGC1333 LH 95 IC 348 Trapezium cluster (Orion Nebula Cluster) NGC 3603 R136 (30 Doradus) Stellar clusters and associations • Stars form in groupings (clusters and associations): 1. Stellar aggregates contain young stars • • These stars are young because of their photospheric characteristics: • Brighter than main sequence – larger radii - still contracting • Show photospheric signatures of infalling gas – H recombination lines • Show evidence of warm dust surrounding them – IR excess from heated dusty disks • Show prominent X-ray emission – typical of young convective stars These stars must have formed recently because they are surrounded by gas 2. Massive stars are found in stellar aggregates • • Massive stars are short lived (few Myr for M>20Msun) – they must have formed recently Therefore the aggregates themselves must be young • Not all these groupings survive as bound clusters • Early dissolution of the stellar population into the field ``infant mortality’’, or slow expansion and stellar loss. • This can be concluded from the comparison of the current star formation rate (stars per unit of time per unit of volume) with the number of stars in open clusters per interval of age. • 90% of the current star local star formation output feeds the galactic field population Trapezium cluster (Orion Nebula Cluster) Pleiades (100Myr) Embedded clusters • At early stages 𝑡𝑡~ < 1 𝑀𝑀𝑀𝑀𝑀𝑀 young clusters are deeply embedded within their parental molecular cloud. • The high dust extinction produced by the high column density of dust limits the detection of their stellar population at optical wavelengths. Absolute magnitude: + Sun: 𝑀𝑀𝑉𝑉 ~5 𝑚𝑚𝑚𝑚𝑚𝑚 Distance modulus: Dust extinction 𝜇𝜇 = 5 log 𝑑𝑑 − 5 + 10pc 𝜇𝜇 = 0 100pc 𝜇𝜇 = 5 1kpc 𝜇𝜇 = 10 10kpc 𝜇𝜇 = 10 Low: 𝐴𝐴𝑉𝑉 < 5 Intermediate 𝐴𝐴𝑉𝑉 ~5 − 15 High 𝐴𝐴𝑉𝑉 ~20 − 100 Trapezium cluster (Orion Nebula Cluster) Approximate limit of optical telescopes: V=20-25 mag (ground based); 28-30 (HST) Embedded clusters • Dust extinction is significantly lower at longer wavelengths. ONC optical Trapezium cluster (Orion Nebula Cluster) • Note: at large distances within the disk of the MW, the buildup of foreground extinction fro the disk itself becomes important too. ONC JH Detection of embedded clusters, NIR surveys Near infrared photometric surveys enable the discovery of embedded young clusters. • • an excess of density of sources in a given location may be cluster candidate color-magnitude and color-color diagram, compared to adjacent fields, can identify the young population and help estimating its properties (number of stars, average age, …) AV direction ONC optical Trapezium cluster (Orion Nebula Cluster) IR excess sources ONC JH IR excess young stars • • • • Some young stars exhibit an excess of flux at long NIR wavelengths (𝜆𝜆 > 2𝜇𝜇𝜇𝜇), as their NIR colors cannot be reproduced invoking only dust extinction. Once the observed multi-wavelength fluxes are corrected for extinction (e.g., based on multi-band photometry at shorter wavelengths, or e.g., knowing the spectral type of the source from spectroscopic measurements, the spectral energy distribution (SED) of the excess is isolated. The excess has a circumstellar origin, and represents the emission of the dust in the inner part of the circumstellar disks, heated by the central star (𝑇𝑇𝑑𝑑𝑑𝑑𝑑𝑑𝑑𝑑 < 𝑇𝑇𝑠𝑠𝑠𝑠𝑠𝑠𝑠𝑠 ) Since circumstellar disks have short lifetimes (<10Myr) as their material is either photoevaporated, accreted, or condenses in planets, circumstellar disk presence from IR excess is an indicator of youth. 𝝆𝝆 Ophiuchi dark cloud complex ONC optical Trapezium cluster (Orion Nebula Cluster) Upper Sco OB association IR spectral index 𝛼𝛼𝐼𝐼𝐼𝐼 = ONC JH 𝑑𝑑 log 𝜆𝜆𝐹𝐹𝜆𝜆 𝑑𝑑 log 𝜆𝜆 IR excess young stars Class I 𝛼𝛼2−10𝜇𝜇𝜇𝜇 > 0 Protostar: emission originates from accreting disk + envelope Class II −1.5 < 𝛼𝛼2−10𝜇𝜇𝜇𝜇 < 0 Pre-main sequence star: IR emission from dusty circumstellar disk ONC optical Trapezium cluster (Orion Nebula Cluster) Class III 𝛼𝛼2−10𝜇𝜇𝜇𝜇 < −1.5 Older pre-main sequence star. Disk partially cleared inside-out ONC JH T and OB-associations • Over a time of ~1Myr embedded clusters start to be optically revealed. We can distinguish different categories. • T-associations: • Systems that were born in low-density dark clouds that did not contain massive stars. • Loose distributions of young pre-main sequence stars • Typical example: the Taurus star forming region. • OB-associations: • The unbound, expanding grouping of newly formed stars, containing massive, bright OB stars. • The high effective temperature of OB stars (𝑇𝑇𝑒𝑒𝑒𝑒𝑒𝑒 > 10,000𝐾𝐾) ionizes the remaining gas creating HII regions • More massive than T associations. • Actual clusters: • When the stellar population remains bound (𝐸𝐸 = 𝐾𝐾𝐾𝐾 + 𝑃𝑃𝑃𝑃 < 0) after gas removal. • Note that a fraction (large or small) of stars may escape the stellar system in its early phases of evolution even when a bound clusters survives afterwards. Taurus association • • • • • Distance: 140 pc Low density of stars: 1-10 stars/pc3 spans over 15 degrees on the plane of the sky Absence of massive stars. Many stars are distributed in relatively compact clumps which follow the density of gas ~1 Myr old (with age spread largely unknown) Luhman (2009) Memberships of young associations Especially (but not only) for the case of loose T associations, the isolation of the young members from field (background and foreground) galactic field stars might be challenging. There are several memberships indicators that can be used; they are in many ways complementary, as each of them alone does not guarantee to accurately identify (if at all) the young nature of PMS stars. Limiting to low-mass PMS stars, these are: 1. Flux excesses and other spectral characteristics • • • Disk excess: as it traces the presence of a circumstellar disk. Note that for more evolved class III sources, the signal might be weak, requiring observations at long wavelengths (MIR, FIR, sub mm) Accretion signatures: as matter falls from the inner disk on to the star, the loss of gravitational potential energy is converted into luminosity, with a typical spectrum including UV excess, optical continuum, and Hydrogen recombination lines. The Hα line at λ=6563 Angstrom excess, in particular, is most commonly used since practical to measure, and bright. A standard classification is as follows: Classical T-Tauri stars: E. W. H𝛼𝛼 > 20 𝐴𝐴 Weak-line T-Tauri stars: 0 < E. W. H𝛼𝛼 < 20 𝐴𝐴 Other accretion indicators are H and K-lines of Ca II at 3968 and 3934 Angstrom X-rays: young low mass stars are generally X-ray emitters, originating in plasma explosively heated and confined in magnetic loops following magnetic reconnection events. This is an analogous of solar flares, and is facilitated by the turbulent and highly convective interiors of low mass PMS stars 2. Kinematic common motions The velocity dispersion of the members of a young star forming region is generally of a few km/s, whereas the dispersion of field stars may be more than one order of magnitude higher. Proper motions and radial velocity surveys help to isolate the candidate members that share the same bulk motion. proper motions of the members in Taurus OB associations An aggregate of young stars containing early (OB) type stars. B-type stars: 𝑇𝑇𝑒𝑒𝑒𝑒𝑒𝑒 > 10,000 𝐾𝐾; 𝑀𝑀 ≳ 4𝑀𝑀⨀ O-type stars: 𝑇𝑇𝑒𝑒𝑒𝑒𝑒𝑒 > 30,000 𝐾𝐾; 𝑀𝑀 ≳ 20𝑀𝑀⨀ The large mass of these stars requires the stellar population to be more numerous than T-associations; thus OB associations originate from more massive clouds. The UV radiation of early type stars is ionizes the surrounding gas, which then emits Hydrogen lines as it recombines forming an HII region. LH 95 (Large Magellanic Cloud) Note that massive star early evolution is fast enough that these objects are already on the main sequence as soon as the stellar system is non embedded. Cygnus OB2 – NGC6914 Spectroscopic analysis can identify if an early type star is in main-sequence or already leaving the main sequence. In the first case, it follows a given temperature-luminosity relation, which can be exploited to determine their distance (but reddening must be corrected for). Nearby young associations and Gould’s belt T associations Many close star forming regions, as well as many of the nearby bright stars, lie on a ring-like structure called Gould’s belt. It has a radius of ~300 pc, and is tilted by 20° with respect to the Galactic plane. Kinematic measurements show that the belt is expanding and slightly rotating. Age ~ 25Myr OB associations Expansion and subgrouping of complexes of star forming regions • • • • • Young star forming regions are often part of larger complexes including subgroups. An example is shown below: the Scorpius-Centaurus association. These subgroups generally formed as part of a joint structure of clouds, or molecular cloud complex. However, different subgroups may have formed at different times Δ𝑡𝑡 ≲ 10 Myr. As young association tend to expand over time, as a result of excess kinetic vs. potential energy, the concentration of an association is a (very rough) indicator of its age. Proper motions in this case can be used to track back the moment in the past when the stellar population was at its maximum density. Upper Scorpius The Orion complex Trapezium cluster Orion Nebula Blue: optical Red: HII Green: H2 Extinction map (Lombardi+2007) Blue: optical Red: HII Green: H2 Orion OB association groups Orion OB Ia: ~12 Myr Orion OB Ib: ~8 Myr Orion OB Ic: ~3 − 6 Myr Orion OB Id: ~2 Myr - Orion Nebula Blue: optical Red: HII Green: H2 Megeath+ (2012) Young low mass stars in the Orion complex, part of the same population of the youngest Orion OB Id population, are identified through their NIR and MIR excess emission from circumstellar material. They are confined in filamentary structures which follow the Orion A and B clouds. Subgrouping on smaller scales in the Orion complex Alves+ (2012) Analysis of the properties of low-mass pre-main sequence stars: lower AV and higher fraction of Class III sources vs Class II to the south of the Orion Nebula may indicate an older population. Ages of young stellar associations 1. Embeddedness (high AV = young). 2. IR excess (temporal evolution Class I → Class II → Class III). 3. 4. Proper motions to track back the expansion of a stellar system Presence of unevolved massive stars sets upper limits to cluster age 5. Turn off age: the luminosity (and mass) of the turn off point from the main sequence decreases with time 6. Pre-main sequence H-R diagram vs evolutionary models 𝑡𝑡𝑀𝑀𝑀𝑀 ∝ 𝑀𝑀−2.5 Pre-main sequence ages • The long contraction time needed for PMS stars to reach the MS can be used as a clock to measure stellar and cluster ages. • BUT: 1) theoretical models for PMS star are more uncertain than those for MS stars 2) determination of 𝑇𝑇𝑒𝑒𝑒𝑒𝑒𝑒 and log 𝐿𝐿 for PMS stars is challenging Embedded massive stars • Massive stars at very young ages can be still deeply embedded • Their UV radiation can create Ultra Compact HII regions • • Smaller than HII regions (𝑟𝑟 < 0.3𝑝𝑝𝑝𝑝) Radiation absorbed and re-emitted at longer wavelengths (MIR, FIR, Radio) Runaway massive stars • 25% of O-type stars are not spatially associated with young clusters or clouds. • Many are located farther from the Galactic plane, and show velocities 𝑣𝑣 = 50 − 100 km/s • Their velocity originates from: • Multi-body stellar encounters • As ex-companions of supernovae Open clusters • The remain of OB associations that have not undergone complete disruption. Size: 2-10pc. • Despite their survival (ages 30-1000Myr) not all of them are bound. • No or little molecular gas left. No active star formation. The members are coeval. • ~1200 known systems (but sample is limited by galactic disk extinction), distributed. • Precious laboratories of stellar evolution: simple stellar populations over a range of ages. Accurate photometric measurements allow to test theoretical models of stellar evolution and spectra of stars. • Draper (1930): discovery of dust extinction: smaller (= farther away) young clusters appear dimmer. Pleiades (130 Myr) NGC4755 (16Myr) Hyades (650Myr) Hyades • The proper motion are aligned towards a convergence point, where the open cluster appear to shrink towards. • This indicates that the cluster is receding. The determination of the convergence point, together with measured proper motions and radial velocities can be used to determine the distance of the cluster: θ 𝑉𝑉𝜏𝜏 𝑉𝑉𝑟𝑟 = tan 𝜃𝜃 ; 𝑉𝑉𝜏𝜏 = 𝜇𝜇 � 𝑑𝑑 𝑉𝑉𝑟𝑟 tan 𝜃𝜃 𝑑𝑑 = 𝜇𝜇 The initial mass function • The IMF is the distribution of the initial masses of stars. 𝑑𝑑𝑑𝑑 𝜉𝜉 𝑚𝑚 = 𝑑𝑑𝑑𝑑 1. Provides useful clues on the processes that lead to the conversion of molecular clouds into stars. 2. Stars of different mass have different lifetimes; the IMF dictates the fate of stellar systems up to entire galaxies. 3. Massive stars provide energy feedback and chemical enrichment. How many does a galaxy form? 4. Very low mass stars and brown dwarfs are ``invisible’’ mass at large distances. How many are there? 5. Its form, universality or variations, is fundamental to interpret observations at large distances (e.g., galaxies) The initial mass function • The IMF is the distribution of the initial masses of stars. • Empirical distribution, based on counting stars of different mass. • The stellar IMF is commonly fitted by a power-law: 𝑑𝑑𝑑𝑑 𝜉𝜉 𝑚𝑚 = ∝ 𝑚𝑚−𝛼𝛼 𝑑𝑑𝑑𝑑 Φ log 𝑚𝑚 = 𝑑𝑑𝑑𝑑 ∝ 𝑚𝑚−Γ 𝑑𝑑 log 𝑚𝑚 𝛼𝛼 = Γ + 1 • Salpeter (1995): the stellar IMF, above 1Msun is well fitted by a single power law 𝛼𝛼 = 2.35 • Extension of the IMF in the low mass stellar range and brown dwarf regime (M<0.08Msun) shows flattening • Kroupa (2001,2001): the galactic field IMF multi-part power law • Chabrier (2005) IMF: lognormal for low-mass stars and brown dwarfs From luminosity function to IMF 1. 𝑀𝑀𝑉𝑉∗ General luminosity function: Distribution of present-day number of stars per absolute magnitude bin (per unit of volume, 𝑝𝑝𝑝𝑝 −3 ) 𝑑𝑑𝑑𝑑 now 𝑑𝑑𝑀𝑀𝑉𝑉 Needs special care when derived: incompleteness, selection effects, and stars of different mass (=luminosity) extend to different scale height above the galactic midplane. Φ 𝑀𝑀𝑉𝑉 = 2. Bright, 𝒎𝒎 ≳ 𝟎𝟎. 𝟖𝟖𝑴𝑴⨀ The first formed already died Initial luminosity function: Φ 𝑀𝑀𝑉𝑉 = 𝑡𝑡𝑔𝑔𝑔𝑔𝑔𝑔 � Ψ 𝑀𝑀𝑉𝑉 = 𝑡𝑡𝑔𝑔𝑎𝑎𝑎𝑎−𝑡𝑡𝑀𝑀𝑆𝑆 (𝑀𝑀𝑉𝑉 ) 𝑡𝑡𝑔𝑔𝑔𝑔𝑔𝑔 � 0 𝑑𝑑𝑑𝑑 at formation 𝑑𝑑𝑀𝑀𝑉𝑉 𝑑𝑑𝑑𝑑 𝑚𝑚̇ 𝑡𝑡 Ψ(MV )/2H(M𝑉𝑉 ) 𝑑𝑑𝑑𝑑 𝑚𝑚̇ 𝑡𝑡 Ψ(MV )/2H(M𝑉𝑉 ) star formation rate scale height 𝑖𝑖𝑖𝑖 𝑀𝑀𝑉𝑉 < 𝑀𝑀𝑉𝑉∗ 𝑖𝑖𝑖𝑖 𝑀𝑀𝑉𝑉 > 𝑀𝑀𝑉𝑉∗ Faint, 𝒎𝒎 ≲ 𝟎𝟎. 𝟖𝟖𝑴𝑴⨀ Never died m 0 t 1010 From luminosity function to IMF 𝜉𝜉 𝑚𝑚 = 𝑑𝑑𝑑𝑑 𝑑𝑑𝑑𝑑 = 𝑑𝑑𝑑𝑑 𝑑𝑑𝑀𝑀𝑉𝑉 𝑑𝑑𝑀𝑀𝑉𝑉 𝑑𝑑𝑑𝑑 = Φ(𝑀𝑀𝑉𝑉 ) 𝑑𝑑𝑀𝑀𝑉𝑉 𝑑𝑑𝑑𝑑 Mass-magnitude relation for MS stars. • Theoretical: from stellar evolutionary simulations and stellar atmosphere modeling. • Empirical: from multiple stellar systems with orbital parameters well determined. Unresolved stellar multiplicity affects the slope of the measured IMF. This correction is sensitive to the actual fraction of stars in multiple systems, its variation with stellar mass (e.g., massive stars more likely to be in multiple systems than low-mass stars), and the companion mass distribution. “stellar IMF” – corrected for multiplicity “system IMF” – uncorrected IMF from young clusters 1. Advantages: - no evolutionary effects. - one single distance for all members of one cluster. - VLMS and BDs are brighter when they are young. 2. Disadvantages: - estimated masses are more uncertain. (both from observational difficulties and theoretical uncertainties in stellar evolutionary models) Orion Nebula cluster, Da Rio+ 2012 Universality of the IMF • There is no evidence of strong systematic variations of the IMF with initial conditions after the first generation of stars in the universe. • The IMF appears universal over several order of magnitude of cluster density, mass, metallicity. • Local variations of the IFM are statistical or exceptional. Bastian+2010