Star Populations and Star Formation

Transcription

1 Star Populations and Star Formation As comprehensive as Cold Dark Matter Theory is, it cannot tell us much about the leading tracer of dark matter, regular, baryonic matter. Galaxy dynamics can tell a lot about a galaxy s evolution, but many other clues are available: patterns in stellar populations chemical abundance patterns Although stars and the ISM have to follow the laws of physics, it is much more difficult to carry out ab initio computations in these areas than it is for a dynamical model of a galaxy where usually only gravity needs to be taken into account Consequently, observations and empirical data play a large role in our understanding of the composite nature of the regular matter in galaxies. Think of a galaxy as an ecosystem! 1

2 Goals Stellar Populations in Galaxies reconstruct the fossil record of star formation, heavy element formation, and galaxy formation delineate systematic trends in stellar content and evolutionary properties of galaxies vs type, mass, bulge/disk/halo/nucleus, radius, environment reconcile results with lookback studies of high-redshift galaxies understand underlying physical processes that drive star formation, galaxy evolution reconcile results with predictions of theoretical scenarios, models, and numerical simulations Diagnostics: what are stellar populations? age chemical abundances (absolute and differential) stellar orbits 2

3 Topics The Solar neighborhood basics of star formation stellar statistics, luminosity function, IMF disk population: ages, abundances, orbits The Galactic halo and bulge field stars formation scenarios: instantaneous collapse vs hierarchical formation globular clusters and the fossil record Resolved stellar populations in nearby galaxies Magellanic Clouds Local group galaxies Integrated stellar populations in galaxies diagnostics of star formation rates and histories spectral and evolutionary synthesis techniques star formation properties of the Hubble sequence starbursts: physical nature and triggering environmental influences on galaxy evolution 3

4 Star Formation Basics I It is simple to show from stellar theory that stars have formed recently in many galaxies: τ MainSequence 3 (M[M ]) yrs (above correct for M<~10M ) Look at typical MS lifetimes: O star M= 30 M τ = 2-3x10 6 yrs K star M= 0.7 M τ = 3x10 10 yrs So many generations of massive stars may have come and gone since a galaxy formed while all the low mass stars ever formed are still around. -- Note that in a typical old galaxy like a giant elliptical, the main sequence turnoff is at about 0.9 M or G5. 4

5 Star Formation Basics II Look at how stars form in sufficient depth to see what might influence the action on galaxy scales Jet from a circumstellar disk returning energy to a cloud Deeply embedded protostar Circumstellar disk Globule before a protostar forms B,V,I Agglomeration & planetesimals Mature planetary system 5

6 Jean s Mass and Cloud Collapse Gravitational force must overcome gas pressure for a cloud to collapse: 2 GM P = M,R = mass,radius of cloud 4 4 π R 2 GM R computed from ave distance between particles nkt= 4 3M 3 n=particle density, m=mass of a particle 4π 4πnm M k T = 2π Gm n Jeans T 2 Jeans = 1 2 M 18 M for a molecular hydrogen cloud, n in particles/cm n Cloud must be very cold and very dense for low mass stars to form, easier to form high mass stars. Magnetic fields and turbulence can hinder cloud collapse. 3 6

7 Modes of Star Formation I: Sequential SF / Self-Propagating SF What causes the first star formation? Many regions in the Milky Way show evidence for previous star formation to have triggered subsequent star formation (Orion is a good example) Clouds can be compressed by radiation pressure or SN shocks or shocks/pressure gradients from expanding HII regions Spiral patterns caused by rotational shear Does not predict a clear relationship between properties across a spiral arm 7

8 Modes of Star Formation II: Density Wave Star Formation But it has been hard to find direct evidence of density-wave triggered star formation -- M51 best studied case. Look at the sequence implied: M99 1. Leading edge of density wave will have compressed molecular gas clouds 2. Immediately behind will be a region with very young stars 3. Further behind older stars will be found Gonzalez & Graham 1996 ApJ 8

9 M51 I: Visible-Hα Imaging Scoville et al

10 M51 II: CO Imaging CO contours superposed on visible light image of M51 (Aalto et al. 1999). Spatially resolved velocities revealed streaming motions expected from spiral density waves although velocities are very high. 10

11 Grand Design versus Flocculent Spirals Flocculent <=> propagating star formation Grand design <=> spiral density wave-induced SF M74 - a grand design spiral NGC a flocculent spiral. 11

12 Modes of Star Formation III: Bimodal Star Formation In the Milky Way, regions of high and low mass star formation are frequently distinct compare the Taurus cloud (no OB stars) with Orion (lots of OB stars) A number of observations suggest that a larger proportion of high mass stars formed early in the Universe Could account for the G Dwarf problem in the Milky Way = why are there so few low metallicity G dwarfs? Could explain why intracluster gas so metal-rich Some extreme starburst galaxies show suppressed low mass star formation Pop III stars almost surely were preferentially very high mass stars (difficult for pure hydrogen to cool) Recent theoretical work has shown that variations in magnetic fields and turbulence inside molecular clouds could lead to bimodal star formation 12

13 Factors Controlling Star Formation Need raw material star formation rate will depend on the amount of gas Need gas agglomerated into clouds clumps of gas need to cool to collapse => easier with more heavy elements (higher metallicity) Triggers to help clouds collapse may be needed to overcome magnetic field pressure possibilities include nearby hot stars or supernovae (sequential star formation) spiral density waves gravitational instabilities in spiral disks but need to avoid situtations where differential rotation (shear) would tear clouds apart 13

14 One More Set of Ingredients: Stellar Evolution Stellar evolutionary time scales provide ages for clusters in the Milky Way and also useful for galaxies Later we ll see that caution is needed when looking at the integrated light of galaxies Consider what colors a single age stellar population would have Need to know the relative numbers of stars with mass = Initial Mass Function 14

15 Schmidt Law: Connecting Gas to a Global SFR How to bridge the gap from what is seen in individual clouds in the MW, for example, to galaxy-wide properties such as the star formation rate (SFR) and stellar birth history? How does the SFR vary from galaxy to galaxy? What factors control the SFR: This is the key we are really looking for since understanding the physics of the SFR would unravel a lot of galaxy evolution. First step in by Schmidt ApJ 1959: SFR ρ n=2 n gas for Schmidt's original work Kennicutt 1989 recast this relation in terms of surface density since that is the observable in other galaxies so N SFR = aσgas where N is used to distinguish the surface density index from the volume density index. How can this relation be measured? Is n=2 universal? How does the constant of proportionality vary? 15

16 Deriving the Schmidt Law Need to measure the current star formation rate and the current gas densities; data on the history of these quantities also highly desirable What are the observables: (many of these most useful for history) Number of stars currently on the main sequence Number of giant branch stars (not very useful because the giant phase is relatively short) Number of white dwarfs Amounts of atomic and molecular gas Metallicity Very young star clusters Other information that is available: 1. stellar evolution rates; 2. fractions of gas returned to the ISM as a result of stellar evolution; 3. chemical yields from SN and other stellar evolution Many processes lock up material that won t or can t be recycled into stars (eg., formation of stars with M<~ 0.9 M, white dwarfs) so gas density must declining with time 16

17 SFR and Gas Density Data I Schmidt s study looked at primarily the solar neighborhood. Due to technology limitations, he could not rely on emission line surveys and had to deduce the SFR by analyzing stellar luminosity functions difficult to execute but yielded a formalism that is useful for looking at the SFR history Kennicutt 1989 executed an improved study of the relationship between SFR and gas density: Hα used as a surrogate for the SFR (but note extinction may be a problem) HI, CO used to measure the gas density (conversion of CO mass to H 2 mass problematic and by necessity assumed to be a constant) Radial profiles of Hα, HI, and CO were available for 15 galaxies ranging from S0 to Sc HI and CO data also included velocities for derivation of rotation curves. 17

18 SFR and Gas Density Data II Lines are for N=1 and N=2 Kennicutt found that the Hα surface brightness correlates most strongly with HI surface mass density, not with the molecular surface mass density. Slope is N~1.4 But local star formation studies obviously show that stars form from molecular clouds, not from HI clouds. CO data sets show excellent correlation between CO mass and SFR. 18

19 SFR and Gas Density Data III From Young and Scoville 1991 Kennicutt investigated a number of possibilities to resolve the discrepancy between the CO results and his Hα results. For example, as shown above he found good correlation between stellar scale lengths and Hα but much poorer between Hα and CO. 19

20 What Gives? Looking at integrated indicators of properties may just show that bigger galaxies have more of everything Plot at left suggests that if the gas density falls below a threshold, star formation is suppressed. 20

21 Martin and Kennicutt 2001: Improvements 21

22 Density Thresholds Martin and Kennicutt 2001 confirmed the earlier results with much higher quality data and a larger sample. Hα 22

23 More on Thresholds NGC 5236 Gas density Critical density Hα gas Toomre Q parameter for stellar disk: Q(R) = σκ π Gµ σ=velocity dispersion κ=epicyclic freq µ= gas surface density 2 σκ µ gas (R) V V dv µ crit = αq α(r)= κ (R)= π G µ crit (R) R R dr If gas density is higher than the critical density, then gravitational instabilities will cause clouds to collapse and stars to form. The velocity dispersion in a disk does not vary much with R while κ varies as 1/R (see Sept 13 lec. notes). => Surface density drops faster than 1/R so it drops below the critical density at some R so this may explain the edges of stellar disks. 23

24 Physical Insight into Q By considering a local region in a disk, one can show that Toomre s Q criterion is the disk analogue to Jean s criterion for cloud collapse Rotation, velocity dispersion provide pressure-like terms that can counteract the gravitational force See B&T p and p Other interesting results include Stability criteria for purely gaseous disks are close to those for stellar disks Stability is only a weak function of disk thickness 24

25 Subcritical Disks shear µ crit π 2.5Aσ dω A=-0.5R G dr Martin and Kennicutt also noted that some galaxies are forming stars in regions where the disk is definitely subcritical: NGC 2403 is a clear example. Rather than the Coriolis force being the key to whether a cloud can collapse, if the shear is small enough, clouds may have time to collapse so a shear criterion is needed in addition to a rotation criterion. 25

26 More on the SFR: What s its History? Collection of stars, stellar remnants, and enriched gas represent the integrated history of star formation in a galaxy Low mass galaxies may lose enriched gas due to SN t(m V ) = Maximum possible age for a star still on the MS ϕ (M V ) = Observed luminosity function for stars on the MS ψ(m V ) = Luminosity function for all stars ever formed N(M V, t) = number of stars in MV 0.5 to MV on MS formed up to time t Assume ψ(m V ) is independent of time, t=0 time of galaxy formation, t=1 is now dn(m V, t) =ψ(m V ) f(t) f(t) = star formation rate dt 1 ϕ (M V) =ψ(m V) f (t)dt 1-t(M ) V ϕ(m v ) is observable, ψ(m V ) can be estimated from open clusters (its equivalent to the initial mass function) so f(t) can be estimated. 26

27 The Solar Neighborhood begin with fundamental stellar statistics Hipparcos parallax survey provides database for >20000 stars out to ~100 pc key statistical properties are luminosity, mass functions key evolutionary diagnostics are ages, kinematics, abundances cf. Kovalevsky 1998, ARAA, 36, 99 Perryman et al. 1997, The Hipparcos Catalog, and A&A, 323, L49 27

28 Luminosity Function LF = number of stars per unit volume per magnitude bin (for faintest stars IR magnitudes needed) Reid, Gizis, Hawley 2002, AJ, 124,

29 Mass-Luminosity Relations converts luminosity function to present day mass function (PDMF) Maeder, Meynet 1988, A&AS, 76,

30 Solar Neighborhood Mass Function Reid et al