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									            Archaeology of Galaxy Evolution in
                     the Local Group

                           Ken Freeman
           Research School of Astronomy & Astrophysics

Evolution of galaxies, their central black holes and their large-scale environment
                                Postdam Sept 2010
       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.
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.
The thin disk: formation and evolution

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 are the stellar velocity dispersions evolve with time

 Measuring stellar ages is still a major problem
The galactic disk shows an abundance gradient
(eg galactic cepheids: Luck et al 2006) ....
+ 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

Metallicity gradient in outer regions of M31disk also bottoms out,
as in the Milky Way
                                                     Worthey et al 2004
The age-metallicity relation in the solar neighborhood is still uncertain

                                                         et al 2006

   The large scatter in                                  Estimating ages
   [Fe/H] at all ages was                                for field stars is
   part of the reason to                                     difficult
   invoke largescale
   radial mixing : bring
   stars from inner and                                Edvardsson et al 1993
   outer Galaxy into the                               Nordstrom et al 2004
   solar neighborhood                                  Valenti & Fisher 2005

                                                              (Reid et al 07)

Preliminary age-metallicity relation for about 400 nearby                    subgiants.

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
Galactic archaeology

                       Radial mixing
Simulation of disk formation from cooling gas in a dark halo potential (Roskar et al 2008)

 Radial mixing: transient spiral arm interactions can move stars from one
 near- circular orbit into another (Sellwood & Binney 2002; Minchev & Famaey 2010).
 How important is this effect in real galaxies ?
     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                                                     W is in the
dispersion    total                                          vertical (z)
  (km/s)                                                      direction

              W = 0.4total
                                                            relation (AVR)
                               stellar age

   (McCormick dwarfs, CaII emission ages)                           Wielen 1977
 • 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.
See also Anguiano et al poster on AVR for                subgiants.

Difficulty of measuring stellar ages is reason for the different views.
                       old disk

                                                Velocity dispersions
                                                of nearby F stars

                                             thick      appears at
                                             disk       age ~ 10 Gyr

                                  Disk heating saturates at 2-3 Gyr

Freeman 1991; Edvardsson et al 1993; Quillen & Garnett 2000
          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
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.

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
LMC     Pompeia, Hill et al. 2008
Sgr     Sbordone et al. 2007
Fornax  Letarte PhD 2007
                                    Abundance ratios reflect different
SculptorHill et al. 2008 in prep        star formation histories
        + Geisler et al. 2005
Carina Koch et al. 2008
        + Shetrone et al. 2003
Milky-Way Venn et al. 2004
                                           •   Each galaxy has had a different
                                               evolutionary track
                                           •   The position of the knee forms a
                                               sequence following SFH-timescales (and
                                               somewhat the galaxy total luminosity)
                                           •    s- process (AGB product) very efficient in
                SNII        +SNIa
                                               galaxies with strong SFR at younger ages
                                               (<5Gyrs): Fnx > LMC > Sgr > Scl
                                           •   r/s-process elements can be used as
                                               another clock (or even 2 clocks: r/s
                                               transition knee, and start of rise in s )
                                           •   AGB lifetimes + s-process yields are
                                               metallicity-dependent (seeds)

                                           •   Abundance pattern in the metal-poor stars
                                               everywhere undistinguishable ? Seems to
                                               be the case for stars in the exended low-
                                               metallicity populations.

         rise in s-process
                                                   Venn 2008
    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

  We are planning a large chemical tagging survey of about a
     million stars, using the new HERMES multi-object
                   spectrometer on the AAT.

The goal is to reconstruct the dispersed star-forming aggregates
that built up the disk, thick disk and halo within about 5 kpc of
                              the sun
          HERMES is a new high resolution multi-
             object spectrometer on the AAT
                    spectral resolution 28,000
                    (high resolution option 50,000)
                    400 fibres over  square degrees
                    4 VPH gratings, 4 bands ~ 1000 Å
                    First light 2012 on AAT

                The four wavelength
             bands are chosen to detect
              lines of elements needed
                for chemical tagging

             Strong synergy with Gaia:
accurate parallaxes (~ 1% errors) and proper motions
    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.
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
                       old disk

                                                Velocity dispersions
                                                of nearby F stars

                                            thick       appears at
                                            disk        age ~ 10 Gyr

                                  Thick disk is kinematically distinct

Freeman 1991; Edvardsson et al 1993; Quillen & Garnett 2000
[( + 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)
                                       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
                                       (eg Gilmore et al 1995)

-2.0   -1.5        -1.0   -0.5   0.0
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.
           • 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

                 Baryonic mass ratio: thick disk/thin disk

                                                             Yoachim & Dalcanton 2006
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.
• 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

            Clump cluster galaxy at z = 1.6
                (Bournand et al 2008)

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.
• 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
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)
                       by massive satellite

                                              Wilson et al 2010, Ruchti
                                              et al 2010: f(e) for thick
                                              disk stars from RAVE :
                                              may favor gas-rich merger
                                              picture ?

 Distribution of orbital eccentricity of
   thick disk stars predicted by the
    different formation scenarios.
                                                                Sales et al 2009
     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

• 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.
The Galactic Stellar Halo
                                                              disk &
                                                            thick disk

                                          |Zmax| < 2 kpc

Rotational velocity of nearby stars relative to the sun vs [m/H]
(V = -232 km/s corresponds to zero angular momentum)
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
                                                                ELS 1986
now - eg Sgr dwarf, NGC 5907

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

     • 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
            <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.
       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.
Stellar halos in other Local Group galaxies - M31, M33, LMC

M33: RR Lyrae stars show an old disk + halo structure (Sarajedini et al 2006)

M31: Spectra of red giants in the outer regions of M31 identify a non-rotating
metal-poor halo (Kalirai et al 2006, Chapman et al 2006) with <[Fe/H]> ~ -1.4

                                                         Giants show an
                                                     abundance gradient out to
                                                     > 100 kpc (bulge + halo)

                                                   Tanaka et al (2009): M31 halo
                                                   shows at least 16 substructures,
                                                   so at least 16 events contributed
                                                   to building halo. Each event has
                                                   mass 107 to 109 M and patchy
                                                   metallicity distribution : ie not
                                                   yet fully mixed.

                                                                 Richardson et al 2009
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 la 2008, Gilbert et al 09,
Cooper et al 2009)
                                   Although the LMC globular clusters do
                                   not lie in a halo, the LMC does have a
                                   kinematically detected stellar halo
                                  The RR Lyrae stars in the LMC have a
                                  velocity dispersion of 53 km/s. This is the
                                  expected dispersion for the LMC's observed
                                  circular velocity. The mean abundance of
                                  these stars is [Fe/H] = -1.45.
                                  (Minniti et al 2004)

Summary for stellar halos:
• most disk galaxies have a stellar halo

• stellar halo is made up mainly of debris of small accreted galaxies,
  although there may be an inner component which formed dissipatively
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.
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          QuickTime™ an d a
                                     FLIC ŽØª° de compressor
unstable to vertical buckling     are need ed to see this picture.

which generates the boxy

This kind of bar/bulge is not
generated by mergers

Shen, 2010
   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) ?
The kinematics of
 the model are as
   observed for
   boxy bulges:
cylindrical rotation

         b = 0.5°
         b = 9.5°
     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 ~ 11,000

              Melissa Ness
            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. Now they lie in the central
bulge region of the Galaxy.
Distributions of the first stars and the metal-free stars

                                                        Brook et al 2007
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
                           The M31
                           Larger velocity dispersion
                           than the Milky Way bulge
                           (140 km s -1)                    J

                                                                  2MASS J image

M31 has a classical bulge plus
   an inner boxy bar/bulge,
  from detailed comparison
of the isophote structure with
       N-body models.

                                           Athanassoula et al (2006), Beaton et al (2007)

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