Structure and Evolution of the Milky Way Ken Freeman Research School of Astronomy & Astrophysics

1. Structure and Evolution of the Milky Way Ken Freeman Research School of Astronomy & Astrophysics RED GIANTS AS PROBES OF THE STRUCTURE AND EVOLUTION OF THE MILKY WAY Rome November 2010

2. The thin disk: formation and evolution Issues: Building the thin disk : its exponential radial structure, the role of mergers Star formation history, chemical evolution, continuing gas accretion Evolutionary processes: disk heating, radial mixing The outer disk: chemical gradient and chemical properties Many of the basic observational constraints are still uncertain: • The star formation history of the disk • How does the metallicity distribution in the disk evolve with time • How do the stellar velocity dispersions evolve with time Measuring stellar ages is still a major problem

3. The galactic disk shows an abundance gradient (eg galactic cepheids: Luck et al 2006) ....

4. + cepheids, other symbols are open clusters in the Galaxy. Clusters have ages 1-5 Gyr, cepheids are younger The abundance gradient and [/Fe]-gradient in the disk has flattened with time, tending towards solar values. For R > 12 kpc, abundance gradient disappears Carney & Yong 2005

5. M31 Metallicity gradient in outer regions of M31disk also bottoms out, as in the Milky Way Worthey et al 2004

6. The age-metallicity relation in the solar neighborhood is still uncertain Rocha-Pinto et al 2006 Estimating ages for field stars is difficult The large scatter in [Fe/H] at all ages was part of the reason to invoke largescale radial mixing : bring stars from inner and outer Galaxy into the solar neighborhood Edvardsson et al 1993 Nordstrom et al 2004 Valenti & Fisher 2005 (Reid et al 07)

7. log age subgiants (Gyr) Our preliminary age-metallicity relation for about 400 nearby subgiants. Ages derived from isochrones in the log g - Te plane via high resolution spectra Gently declining A-M relation with rms scatter of only 0.15 dex in [M/H] (scatter includes the [M/H] error of ~ 0.10). Less need for radial mixing. Wylie de Boer et al 2010

8. What is the observed form of the heating with time ? The observational situation is not yet clear ... • One view is that stellar velocity dispersion  ~ t 0.2-0.5 eg Wielen 1977, Dehnen & Binney 1998, Binney et al 2000 … velocity dispersion (km/s) W is in the vertical (z) direction total Wielen age-velocity relation (AVR) W = 0.4total stellar age (McCormick dwarfs, CaII emission ages) Wielen 1977

9. • Another view is that heating occurs for the first ~ 2 Gyr, then saturates because stars are mostly away from the Galactic plane Edvardsson et al (1993) measured accurate individual velocities and ages for ~ 200 subgiants near the sun. Their data indicate heating for the first ~ 2 Gyr, with no significant subsequent heating. Disk heating in the solar neighborhood appears to saturate after 2 Gyr, when z ~ 20 km/s. Soubiran et al (2008) measured sample of clump giants, and agree. Difficulty of measuring stellar ages is reason for the different views. Accurate ages from asteroseismology would be very welcome. Accurate ages and distances for giants would allow us to measure the AMR and AVR out to several kpc from the sun.

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

11. The Formation of the Thick Disk Thick disk Most spirals (including our Galaxy) have a second thicker disk component The thick disk and halo of NGC 891 (Mouhcine et al 2010): thick disk has scale height ~ 1.4 kpc and scalelength 4.8 kpc, much as in our Galaxy.

12. Our Galaxy has a significant thick disk • its scaleheight is about 1000 pc, compared to 300 pc for the thin disk • its surface brightness is about 10% of the thin disk’s. • it rotates almost as rapidly as the thin disk • its stars are older than 10 Gyr, and are • significantly more metal poor than the thin disk : mostly(-0.5 > [Fe/H] > -1.0) • alpha-enriched so its star formation was rapid From its kinematics and chemical properties, the thick disk appears to be a discrete component, distinct from the thin disk

13. appears at age ~ 10 Gyr thick disk old disk Velocity dispersions of nearby F stars Thick disk is kinematically distinct Freeman 1991; Edvardsson et al 1993; Quillen & Garnett 2000

14. [( + Eu)/H] vs [Fe/H] for thin and thick disks near the sun The thick disk is chemically distinct Navarro et al (2010), Furhmann (2008)

15. thin disk 225 pc thick disk 1048 pc halo Veltz et al (2008) analysed the kinematics of stars near the Galactic poles in terms of components of different W. The figure shows the weights of the components: the kinematically distinct thin and thick disks and the halo are evident.

16. -2.0 -1.5 -1.0 -0.5 0.0 [Fe/H] Ivezic et al 2008 see the thick disk up to z ~ 4 kpc: [Fe/H] between -0.5 and -1.0 current opinion is that the thick disk itself shows no vertical abundance gradient (eg Gilmore et al 1995)

17. The old thick disk is a very significant component for studying Galaxy formation, because it presents a kinematically and chemically recognizable 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.

18. Baryonic mass ratio: thick disk/thin disk • Most disk galaxies have thick disks • The fraction of baryons in the thick disk is typically small (~ 10-15%) in large galaxies like the MW but rises to ~ 50% in smaller disk systems Yoachim & Dalcanton 2006

19. How do thick disks form ? • a normal part of early disk settling : energetic early star forming events, eg gas-rich merger (Samland et al 2003, Brook et al 2004) • accretion debris (Abadi et al 2003, Walker et al 1996). The accreted galaxies that built up the thick disk of the Galaxy would need to be more massive than the SMC to get the right [Fe/H] abundance (~ - 0.7) The possible discovery of a counter-rotating thick disk (Yoachim & Dalcanton 2008) would favor this mechanism.

20. Clump cluster galaxy at z = 1.6 (Bournand et al 2008) • heating of the early thin disk by disruption of massive clusters (Kroupa 2002). The internal energy of the clusters is enough to thicken the disk Much recent work on the significance of these high-z clump structures as origin of metal-rich globular clusters (Shapiro et al 2010); origin of thick disk and bulges via merging of clumps and heating by clumps (Bournaud et al 2009) Clumps form by gravitational instability, generate thick disks with uniform scale height rather than the flared thick disks generated by minor mergers. (Recall diamond shape of the thick disk of NGC 891)

21. • early thin disk, heated by accretion events - eg the  Cen accretion event (Bekki & KF 2003): Thin disk formation begins early, at z = 2 to 3. Partly disrupted during active merger epoch which heats it into thick disk observed now, The rest of the gas then gradually settles to form the present thin disk • thick disk is generated by radial mixing of more energetic stars from the inner early disk (eg Schönrich & Binney 2009)

22. How to test between these possibilities for thick disk formation ? Sales et al (2009) looked at the expected orbital eccentricity distribution for thick disk stars in different formation scenarios. Their four scenarios are: • a gas-rich merger (Brook et al 2004, 2005). The thick disk stars are born in-situ • accretion (Abadi 2003). The thick disk stars come in from outside • heating of the early thin disk by accretion of a massive satellite • radial migration (stars on more energetic orbits migrate out from the inner galaxy to form a thick disk at larger radii where the potential gradient is weaker (Schönrich & Binney 2009)

23. by massive satellite Abadi (gas-rich) Wilson et al 2010, Ruchti et al 2010: f(e) for thick disk stars from RAVE : may favor gas-rich merger picture ? Firm control of selection effects is needed in separation of thin and thick disk stars Distribution of orbital eccentricity of thick disk stars predicted by the different formation scenarios. Sales et al 2009

24. Thick disk summary • Thick disks are very common. • In our Galaxy, the thick disk is old, and kinematically and chemically distinct from the thin disk. What does it represent in the galaxy formation process ? • The orbital eccentricity distribution will provide some guidance. • Chemical tagging will show if the thick disk formed as a small number of very large aggregates, or if it has a significant contribution from accreted galaxies. This is one of the goals for the HERMES survey.

25. The Galactic Stellar Halo

26. rapidly rotating disk & thick disk slowly rotating halo |Zmax| < 2 kpc Rotational velocity of nearby stars relative to the sun vs [m/H] (V = -232 km/s corresponds to zero angular momentum)

27. Widely believed now that the stellar halo ([Fe/H] < -1) comes mainly from accreted debris of small satellites - cf Searle & Zinn 1978 • Is there a halo component that formed dissipationally early in the Galactic formation process ? eg ELS, Samland & Gerhard 2003 Halo- building accretions are still happening now - eg Sgr dwarf, NGC 5907 ELS 1986 Satellite metallicity distributions not like the metallicity distribution in the halo (Venn 08) - but maybe were more alike long ago. Fainter satellites are more metal-poor and consistent with the MW halo in their [alpha/Fe] behaviour

28. NGC 5907: debris of a small accreted galaxy Our Galaxy has a similar structure from the disrupting Sgr dwarf APOD

29. • Is there a halo component that formed dissipationally early in the Galactic formation process ? Hartwick (1987) : metal-poor RR Lyrae stars show a two-component halo: a flattened inner component and a spherical outer component. Carollo et al (2010 ) identified a two-component halo plus thick disk in sample of 17,000 SDSS stars, mostly with [Fe/H] < -0.5 Describe kinematics well with these three components: <V>  [Fe/H] Thick disk 182 51 -0.7 Inner halo 7 95 -1.6 Outer halo -80 180 -2.2 (retrograde) From comparison with simulations, Zolotov et al (2009) argue that the inner halo has a partly dissipational origin while the outer halo is made up from debris of faint metal-poor accreted satellites.

30. Nissen & Schuster (2010): 78 halo stars - see high and low [alpha/Fe] groups. Abundances [Fe/H] > -1.6 Low [/Fe] stars are in mostly retrograde orbits The high-alpha stars could be ancient in-situ stars, maybe heated by satellite encounters. The low-alpha stars may be accreted from dwarf galaxies. Note different V-distribution of red and blue points.

31. How much of halo comes from accreted structures ? Ibata et al (2009) ACS study of halo of NGC 891 (nearby, like MW, but not Local Group) shows lumpy metallicity distribution, indicating that its halo is made up largely of accreted structures which have not yet mixed away. (cf simulations of stellar halos by Font et al 2008, Gilbert et al 09, Cooper et al 2009) APOD

32. Summary for the Galactic stellar halo: • stellar halo is made up mainly of debris of small accreted galaxies, although there may be an inner component which formed dissipatively

33. The Galactic bar/bulge The boxy appearance of the bulge is typical of galactic bars seen edge-on. Where do these bar/bulges come from ? They are very common. About 2/3 of spiral galaxies show some kind of central bar structure in the infra-red.

34. The bars come naturally from instabilities of the disk. A rotating disk is often unstable to forming a flat bar structure at its center. This flat bar in turn is often unstable to vertical buckling which generates the boxy appearance. This kind of bar/bulge is not generated by mergers

35. The maximum vertical extent of boxy/peanut bulges occurs near radius of vertical and horizontal Lindblad resonances ie where b =  - /2 =  - z/2 (both  and z depend on the amplitude of the oscillation) Stars in this zone oscillate on orbits which support the peanut shape. In-plane End on Edge-on Orbits supporting the peanut

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

37. b = 0.5° b = 9.5° The kinematics of the model are as observed for boxy bulges: cylindrical rotation

38. The stars of the bulge are old and enhanced in -elements  rapid star formation history If the bar formed from the disk, then are the bulge stars and adjacent disk stars chemically similar ? Not clear yet Here the data for the bulge stars and thick disk stars come from different sources [/Fe] higher for thick disk than for thin disk: higher still for bulge Fulbright et al 2007

39. bulge Differential analysis of O-abundance in bulge, thick disk and thin disk stars. The thick disk is O-enhanced relative to thin disk as usual, but the bulge and thick disk look very similar in this study. thick disk thin disk Meléndez et al 2008

40. The bar-forming and bar-buckling process takes 2-3 Gyr to act after the disk settles In the bar-buckling instability scenario, the bulge structure is probably younger than the bulge stars, which were originally part of the inner disk The alpha-enrichment of the bulge and thick disk comes from the rapid chemical evolution which took place in the inner disk before the instability acted The stars of the bulge and adjacent disk should have similar ages in this scenario. Accurate asteroseismology ages for giants of the bulge and inner disk would be a very useful test of the scenario

41. Melissa Ness If the bulge comes from disk instabilities, then the stars in the bulge were once part of the inner disk: its stars are older than the bulge structure We are doing a survey of about 28,000 clump giants in the bulge and the adjacent disk, to measure the chemical properties of stars (Fe, Mg, Ca, Ti, Al, O) in the bulge and adjacent disk: are they similar, as we would expect if the bar/bulge grew out of the disk ? We use the AAOmega fiber spectrometer on the AAT, to acquire medium-resolution spectra of about 350 stars at a time : R ~ 12,000

43. Where are the first stars now ? Diemand et al 2005, Moore et al 2006, Brook et al 2007 … The metal-free (pop III) stars formed until z ~ 4 in chemically isolated sub- halos far away from largest progenitor. If its stars survive, they are spread through the Galactic halo. If they are not found, then their lifetimes are less than a Hubble time  truncated IMF The oldest stars form in the early rare density peaks that lay near the highest density peak of the final system. They are not necessarily the most metal-poor stars in the Galaxy. Now they lie in the central bulge region of the Galaxy. Accurate asteroseismology ages for metal-poor stars in the inner Galaxy would provide a great way to tell if they are the oldest stars or just stars of the inner Galactic halo. Needs ~10% precision in age.

44. Distributions of the first stars and the metal-free stars Brook et al 2007

45. Bulge rotation for metal rich and metal poor stars • Is there a small classical merger-generated bulge component, in addition to the boxy/peanut bar/bulge which probably formed from the disk ? • See a slowly rotating metal-poor component of the bulge. How do we identify the first stars from among the metal-poor stars in the bulge region ? Ness et al 2010

46. The goals of galactic archaeology We seek signatures or fossils from the epoch of Galaxy assembly, to give us insight about the processes that took place as the Galaxy formed. A major goal is to identify observationally how important mergers and accretion events were in building up the Galactic disk and the bulge. CDM predicts a high level of merger activity which conflicts with many observed properties of disk galaxies.

47. Aim to reconstruct the star-forming aggregates and accreted galaxies that built up the disk, bulge and halo of the Galaxy Some of these dispersed aggregates can be still recognised kinematically as stellar moving groups. For others, the dynamical information was lost through heating and mixing processes, but they are still recognizable by their chemical signatures (chemical tagging). Try to find groups of stars, now dispersed, that were associated at birth either • because they were born together in a single Galactic star-forming event, or • because they came from a common accreted galaxy.

48. Stellar substructures in the disk The galactic disk shows kinematical substructure in the solar neighborhood: groups of stars moving together, usually called moving stellar groups (Kapteyn, Eggen) • Some are associated with dynamical resonances (eg Hercules group): don't expect to see chemical homogeneity or age homogeneity (eg Antoja et al 2008, Famaey et al 2008) • Some are debris of star-forming aggregates in the disk (eg HR1614 group and Wolf 630 group). They are chemically homogeneous; such groups could be useful for reconstructing the history of the galactic disk. • Others may be debris of infalling objects, as seen in CDM simulations: eg Abadi et al 2003

49. Look at the HR1614 group (age ~ 2 Gyr, [Fe/H] = +0.2) which appears to be a relic of a dispersed star forming event. Its stars are scattered all around us. This group has not lost its dynamical identity despite its age. De Silva et al (2007) measured accurate differential abundances for many elements in HR1614 stars, and found a very small spread in abundances. This is encouraging for recovering dispersed star forming events by chemical tagging The HR 1614 group is probably the dispersed relic of an old star forming event. V U

50. Chemical studies of the old disk stars in the Galaxy can help to identify disk stars which came in from outside in disrupting satellites, and also those that are the debris of dispersed star-forming aggregates like the HR 1614 group (Freeman & Bland-Hawthorn 2002) The chemical properties of surviving satellites (the dwarf spheroidal galaxies) vary from satellite to satellite, and are different in detail from the more homogeneous overall properties of the disk stars. We can think of a chemical space of abundances of elements O, Na, Mg, Al, Ca, Mn, Fe, Cu, Sr, Ba, Eu for example. The dimensionality of this space is between about 7 and 9. Most disk stars inhabit a sub-region of this space. Stars which came in from satellites may be different enough to stand out from the rest of the disk stars. With this chemical tagging approach, we may also be able to detect or put observational limits on the satellite accretion history of the galactic disk