Slide1: Saas-Fee Lectures 2007 The Origin of the Galaxy and the Local Group Ken Freeman, RSAA, ANU 5. The Galaxy: Bulge, Thick Disk and Thin Disk


Slide2: The Galactic Bar- Bulge small exponential bulge - typical of later-type galaxies. Unlike the large r1/4 bulge of M31 Launhardt 2002 Pritchet & van den Bergh 1994 M31


Slide3: Age and metallicity of the bulge Zoccali et al 2003 : stellar photometry at (l, b) = ( 0º.3, -6º.2) : old population > 10 Gyr. No trace of younger population. Extended metallicity distribution, from [Fe/H] = -1.8 to +0.2 (ie not very metal-rich at |b| = 6º ) Bulge MDF covers similar interval to (thin disk + thick disk) near sun •


Slide4: Inhomogeneous collection of photometric ( ) and spectroscopic ( ) mean abundances - evidence for abundance gradient along minor axis of the bulge Minniti et al 1995 ( kpc ) Abundance gradient in the bulge Zoccali et al (2003)


Slide5: Near the center of the bar/bulge is a younger population, on scale of about 100 pc : the nuclear stellar disk (M ~ 1.5 x 109 M_sun) and nuclear stellar cluster (~ 2 x 107 M_sun ) in central ~ 30 pc. (Launhardt et al 2002) ~ 70% of the luminosity comes from young main sequence stars.


Slide6: The bulge globular clusters Dinescu et al 2003 3D kinematics of 7 globular clusters in the bar/bulge Their velocities show: • all of them are confined to the bulge region • the metal-poor clusters (o) are part of the inner halo • the metal-rich clusters include • a bar cluster • clusters belonging to a rotationally supported system


Slide7: Later type galaxies like the Milky Way mostly have small near-exponential boxy bulges, rather than r1/4 bulges. (eg Courteau et al 1996) These small bulges are probably not merger products: more likely generated by disk instability Boxy bulges, as in our Galaxy, are associated with bars, believed to come from bar-buckling instability of disk. theory: eg Combes & Sanders 1981 ... observations: eg Bureau & KF 1999 ... How did the Galactic Bulge form ?


Slide8: NGC 5746: gas kinematics in a boxy bulge show the signature of orbits in a bar potential (Bureau & Freeman 1999)


Slide9: Our bar-bulge is ~ 3.5 kpc long, axial ratio ~ 1: 0.3: 0.3 pointing about 20-35o from sun-center line into first quadrant (eg Bissantz & Gerhard 2002).


Slide10: López et al (2006) find evidence of a longer flat bar lying in the disk of the Galaxy (7.8 x 1.2 x 0.2 kpc) from 2MASS counts and red-clump stars. The central boxy bar/bulge is the inner extended part of this longer flat bar  GC


Slide11: The stars of the bulge are old and enhanced in -elements  rapid star formation history In the bar-buckling instability scenario, the bulge structure may be younger than its stars, which were originally part of the inner disk


Slide12: The galactic bulge is rotating, like most other bulges: (Kuijken & Rich (2002) HST proper motions) Rotation (Beaulieu et al 2000) K giants from several sources and planetary nebulae (+) Velocity dispersion of inner disk and bulge are fairly similar not easy to separate inner disk and bulge kinematically Bulge ends at |l| ~ 12o


Slide13: exponential in R and z : scaleheight ~ 300 pc, scalelength 3-4 kpc (!) velocity dispersion decreases from ~ 100 km/s near the center (similar to bulge) to ~ 15 km/s at 18 kpc Lewis & KCF 1989 2 1.5 1 R (kpc) log (velocity dispersion) Velocity dispersion of the thin disk


Slide14: Compare the structure and kinematics of the galactic bulge with an N-body simulation of a disk that has generated a boxy bar/bulge through bar-buckling instability of the disk (Athanassoula). Do the simulations match the properties of the Galactic bar/bulge (eg exponential stucture, cylindrical rotation ?) How to test whether the bulge formed through the bar-buckling instability of the inner disk ?


Slide15: N-body model seen from galactic pole


Slide16: COBE Minor axis surface brightness profiles The slope of log I(b) gives the length scale for the model log intensity |b| N-body model


Slide17: The kinematics of the model are as observed for other boxy bulges: eg cylindrical rotation b = 0.5 b = 9.5 Detailed velocity data not yet available for the galactic bar/bulge


Slide18: Vrot l Rotation of bulge (5 < |b| < 10) model V rot (l ) gives the velocity scale for the model (km/s)


Slide19: Velocity dispersion of bulge (5 < |b| < 10) model (km/s)  los l


Slide20: The bulge is not a dominant feature of our Galaxy - only about 25% of the light. The bulge is probably an evolutionary structure of the disk, rather than a feature of galaxy formation in the early universe. Structure and kinematics (so far) are well represented by product of disk instability. The -enhancement indicates that star formation in this inner disk/bulge region proceeded rapidly. The bulge structure may be younger than its stars. The Galactic Bulge - summary


Slide21: NGC 4762 - a disk galaxy with a bright thick disk (Tsikoudi 1980) Most spirals (including our Galaxy) have a second thicker disk component, believed to be the early thin disk heated by an accretion event. In some galaxies, it is easily seen : The thin disk The thick disk The Galactic Thick Disk


Slide22: The Galactic thick disk is detected in star counts. Its larger scale height means its velocity dispersion is higher than for the thin disk and therefore its rotation lags the LSR by more. Near the sun, the galactic thick disk is defined mainly by stars with [Fe/H] in the range -0.5 to -1.0, though its MDF has a tail extending to very low [Fe/H] ~ -2.2. The thick disk appears to be a discrete component, distinct from the thin disk


Slide23: Kinematics and structure of the thick disk rotational lag ~ 30 km/s near the sun (Chiba & Beers 2000) and increases by about 30 km s-1 kpc-1 with height above the plane (Girard et al 2006) velocity dispersion in (U,V,W) = (46, 50, 35) km/s radial scale length = 3.5 to 4.5 kpc : uncertain scale height from star counts = 800 to 1200 pc (thin disk ~ 300 pc) density = 4 to 10% of the local thin disk


Slide24: Freeman 1991; Edvardsson et al 1993; Quillen & Garnett 2000 Velocity dispersions of nearby F stars old disk thick disk Thick disk is discrete component appears at age ~ 10 Gyr


Slide25: The Galactic thick disk is old (> 12 Gyr) and significantly more metal poor than the thin disk: mean [Fe/H] ~ -0.7 and -enhanced  rapid chemical evolution thick disk thin disk higher [/Fe]  more rapid formation


Slide26: [Eu/Fe] (r-process) enhanced by ~ 0.4 [Ba/Fe] deficient by ~ 0.2 (Reddy 06) current opinion is that the thick disk shows no vertical abundance gradient (eg Gilmore et al 1995) [/Fe] enhanced by about 0.25 - many authors Thick disk element ratios


Slide27: Thick disks are very common Almost all spirals have one (Dalcanton & Bernstein 2002) The age distribution for the thick disk stars indicates a time delay between formation of thick disk stars and the onset of star formation in the current thin disk.


Slide28: Dalcanton & Bernstein 2002 Surface photometry (BRK) of 47 late-type edge-on galaxies: Find that all are embedded in a flattened low surface brightness red envelope or thick disk Age > 6 Gyr, not very metal-poor, like thick disk of the Milky Way Formation of the thick disk is a nearly universal feature of formation of disk galaxies


Slide29: The thick disk is a very significant component for studying galaxy formation, because it presents a kinematically recognizable ‘snap-frozen’ relic of the early galaxy. Secular heating is unlikely to affect its dynamics significantly, because its stars spend most of their time away from the galactic plane.


Slide30: Formation pictures for the thick disk ... • a normal part of disk settling (eg Samland et al 2003) • accretion debris (Steinmetz et al 2003, Walker et al 1996) • early thin disk, heated by accretion events - eg the  Cen accretion event (Bekki & KF 2003)


Slide31: The rest of the gas then gradually settles to form the present thin disk My favored formation picture for the galactic disk Thin disk formation begins early, at z = 2 to 3 Partly disrupted during merger epoch which heats it into thick disk observed now Not much is known about the radial extent of the thick disk. This is important, if the thick disk really is the heated early thin disk. Disks form from inside out, so the extent of the thick disk now would reflect the extent of the thin disk at the time of heating.


Slide32: Many of the oldest stars in the disk are debris from accreted satellites which ends up in the thin and thick disk. CDM simulations of formation of an early-type disk galaxy (Abadi et al 2003) show that not all disk stars form in the disk Satellite orbit is dragged into disk plane by dynamical friction - acts like dissipation, although system is collisonless Formation of disk stars outside the disk


Slide33: Why is this interesting ? Because we see thick disk stars in the solar neighborhood with [Fe/H] abundances as low as the most metal-poor globular clusters. Did these stars form as part of early disk formation, or were they acquired ?


Slide34: The thick disk forms rapidly and early (12 Gyr ago in the Galaxy) Appears to be distinct from the thin disk Formed by heating of the early thin disk in an epoch of merging which ended ~ 12 Gyr ago (eg Quinn & Goodman 1986) or from early accretion of satellites, probably in mainly gaseous form (eg Brook et al 2004) Summary of thick disks


Slide35: There is at least one spiral which does not have a thick disk


Slide36: NGC 4244 (MB = - 18.4) appears to be a pure thin disk: just a single exponential component, no thick disk Fry et al 1999


Slide37: The existence of such a pure disk galaxy is interesting because, for at least some late-type disks • the star formation did not start until the gas had settled to the disk plane • since the onset of star formation in the disk, the disk has suffered no significant dynamical disturbance from internal or external sources (pure disk galaxies are not readily produced in CDM simulations: too much merger activity)


Slide38: The Thin Disk


Slide39: star formation history in galactic thin disk : roughly uniform, with episodic star bursts for ages 10 Gyr Rocha-Pinto et al (2000)


Slide40: old disk thick disk No significant chemical evolution in the nearby old disk for ages 2-10 Gyr


Slide41: The outer disk of the Galaxy The galactic disk shows an abundance gradient, as in M31 (eg galactic cepheids: Luck et al 2006 - young stars)


Slide42: Not a simple axisymmetric gradient : Luck et al (2006) cepheids


Slide43: Yong & Carney 2005; Carney & Yong 2005: high resolution spectra of open clusters and stars in the outer disk The abundance gradient for the open clusters (ages 1 to 5 Gyr) bottoms out, at RG = 12 kpc (RG = 15 kpc in M31), and at an abundance of [Fe/H] = -0.5 (as in M31). Outer disk is -enhanced, with [/Fe] = + 0.2 (also Eu-enhanced): indicates fairly rapid star formation history in the outer disk, unlike the solar neighborhood.


Slide44: Solar neighborhood kinematics: Several mechanisms for heating disk stars: transient spiral arms, GMC scattering accretion events, Expect heating mechanisms to saturate after a few Gyr as heated stars spend less time near galactic plane What do observations show ?


Slide45: Freeman 1991; Edvardsson et al 1993; Quillen & Garnett 2000 Velocity dispersions of nearby F stars old disk thick disk Disk heating saturates at 2-3 Gyr appears at age ~ 10 Gyr


Slide46: Nordstrom et al (2004) do not see this saturation at 2-3 Gyr W age(Gyr) velocity dispersion (km/s)


Slide47: Need to resolve this issue If the heating does continue, it will make it more difficult to trace old moving stellar groups which we would like to use to identify substructures in the thin disk.


Slide48: We would like to extend the approach of reconstruction to the disk and thick disk of the Galaxy. Understanding disk formation is more important than understanding the halo, because most of the galactic baryons are in the disk. Disk Reconstruction


Slide49: Like the halo, the galactic disk also shows kinematical substructure : usually called moving stellar groups. Not all of these moving groups are fossils • Some are associated with dynamical resonances (bar) (Hercules group: Dehnen, OLR, new high velocity group: Fuchs, 4:1) • Some are debris of star-forming aggregates in the disk. (HR1614: Feltzing & Holmberg, de Silva et al 2006) • Others may be debris of infalling objects, as seen in CDM simulations: eg Arcturus group : Navarro et al 2004, Mary Williams 2006.


Slide50: We would like to reconstruct the ancient star-forming aggregates of the disk: phase mixing has dispersed them azimuthally right around the Galaxy Structurally invisible - may see them in velocity space (disk moving groups) and integral space (as for halo) and in their chemical properties. May be able to detect evolution of the cluster mass function, the star formation rate, and epochs of satellite infall and star-bursts during the formation of the disk.


Slide51: Use the detailed chemical abundances of thick disk stars ([Fe/H], [/Fe], r- and s- process elements) to tag them to common ancient star-forming aggregates with similar abundance patterns (eg Freeman & Bland-Hawthorn ARAA 2002) The detailed abundance pattern reflects the chemical evolution of the gas from which the aggregate formed. Chemical Tagging Different supernovae provide different yields (depending on mass, metallicity, detonation details, ejected mass ...) leading to scatter in detailed abundances, especially at lower metallicities (enrichment by only a few SN)


Slide52: Chemical tagging is not just assigning stars chemically to a particular population (thin disk, thick disk, halo) Chemical tagging is intended to assign stars chemically to substructure which is no longer detectable kinematically


Slide53: We might expect that groups originating from a common star formation episode have common ages and chemical properties. Chemical Homogeneity of stellar moving groups Look at the HR1614 group (age ~ 2 Gyr, [Fe/H] = +0.2). Studied by Feltzing & Holmberg (2000) who argued for its reality. De Silva (2006) measured precise chemical abundances for many elements in HR1614 stars (AAT: UCLES), and finds a very high level of homogeneity in abundances and in ages. Group is very well defined kinematically despite its age.


Slide54: HR1614 group: [Fe/H] abundances relative to the star HR1614 (Na, Al, Mg, Si, Ca, Mn, Ni, Zr, Ba, Ce, Nd, Eu are similarly homogenous) De Silva 2006


Slide55: HR1614 moving group stars: the (U,V) plane Epicyclic theory predicts constant V. (Eggen, Woolley) The small tilt is expected because epicyclic theory is not valid for these larger V-values. De Silva 2006


Slide56: Conclusion HR1614 is an intermediate age (2 Gyr) moving stellar group which has survived the heating processes in the disk. Its members can be identified kinematically and chemically. This gives us encouragement that disk reconstruction may be feasible using kinematical and chemical identifiers. WFMOS on Subaru/Gemini: 1500 high resolution spectra simultaneously