Annu. Rev. Astron. Astrophys. 2002. 40:487-537
Copyright © 2002 by Annual Reviews. All rights reserved |
Our framework is that the Galaxy formed through hierarchical aggregation. We identify three major epochs in which information about the proto-hierarchy is lost:
We classify signatures relative to these three epochs. The role of the environment is presently difficult to categorize in this way. Environmental influences must be operating across all of our signature classes.
4.1. Zero Order Signatures – Information Preserved Since Dark Matter Virialized
During the virialization phase, a lot of information about the local hierarchy is lost: This era is dominated by merging and violent relaxation. The total dark and baryon mass are probably roughly conserved, as is the angular momentum of the region of the hierarchy that went into the halo. The typical density of the environment is also roughly conserved: Although the structure has evolved through merging and relaxation, a low density environment remains a low density environment (see White 1996 for an overview).
4.1.2. Signatures of the environment
The local density of galaxies (and particularly the number of small satellite systems present at this epoch) affects the incidence of later interactions. For the Local Group, the satellite numbers appear to be lower than expected from CDM (Moore et al. 1999, Klypin et al. 1999). However there is plenty of evidence for past and ongoing accretion of small objects by the Milky Way and M31 (Ibata et al. 1995, 2001b).
The thin disk component of disk galaxies settles dissipatively in the potential of the virialized dark halo (e.g., Fall & Efstathiou 1980). The present morphology of the thin disk depends on the numbers of small galaxies available to be accreted: A very thin disk is an indication of few accretion events (dark or luminous) after the epoch of disk dissipation and star formation (e.g., Freeman 1987, Quinn et al. 1993, Walker et al. 1996). The formation of the thick disk is believed to be associated with a discrete event that occurred very soon after the disk began to settle, at a time when about 10% of the stars of the disk had already formed. In a low density environment, without such events, thick disk formation may not occur. Since the time of thick disk formation, the disk of the Galaxy appears to have been relatively undisturbed by accretion events. This is consistent with the observation that less than 10% of the metal-poor halo comes from recent accretion of star forming satellites (Unavane et al. 1996).
The existence and structure of the metal-poor stellar halo of the Galaxy may depend on accretion of small objects. This accretion probably took place after the gaseous disk had more or less settled – the disk acts as a resonator for the orbit decay of the small objects. So again the environment of our proto-galaxy may have a strong signature in the very existence of the stellar halo, and certainly in its observed substructure. We would not expect to find a stellar halo encompassing pure disk galaxies, consistent with the limited evidence now available (Freeman et al. 1983, Schommer et al. 1992).
4.1.3. Signatures of global quantities
During the process of galaxy formation, some baryons are lost to ram pressure stripping and galactic winds. Most of the remaining baryons become the luminous components of the galaxy. The total angular momentum J of the dark halo may contribute to its shape, which in turn may affect the structure of the disks. For example, warps may be associated with misalignment of the angular momentum of the dark and baryonic components. The dark halo may have a rotating triaxial figure; the effect of a rotating triaxial dark halo on the self-gravitating disk has not yet been seriously investigated (see Bureau et al. 1999).
The binding energy E at the epoch of halo virialization affects the depth of the potential well and hence the characteristic velocities in the galaxy. It also affects the parameter λ = J|E|1/2 G−1 M−5/2, where M is the total mass: λ, which is critical for determining the gross nature of the galactic disk as a high or low surface brightness system (e.g., Dalcanton et al. 1997).
The relation between the specific angular momentum J / M and the total mass M (Fall 1983) of disk galaxies is well reproduced by simulations (Zurek et al. 1988). Until recently, ellipticals and disk galaxies appeared to be segregated in the Fall diagram: From the slow rotation of their inner regions, estimates of the J / M ratios for ellipticals were about 1 dex below those for the spirals. More recent work (e.g., Arnaboldi et al. 1994) shows that much of the angular momentum in ellipticals appears to reside in their outer regions, so ellipticals and spirals do have similar locations in the Fall diagram. Internal redistribution of angular momentum has clearly occurred in the ellipticals (Quinn & Zurek 1988).
The remarkable Tully-Fisher law (1977) is a correlation between the HI line-width and the optical luminosity of disk galaxies. It appears to relate the depth of the potential well and the baryonic mass (McGaugh et al. 2000). Both of these quantities are probably roughly conserved after the halo virializes, so the Tully-Fisher law should be regarded as a zero-order signature of galaxy formation. The likely connecting links between the (dark) potential well and the baryonic mass are (a) a similar baryon/dark matter ratio from galaxy to galaxy, and (b) an observed Faber-Jackson law for the dark halos, of the form M ∝ σ4: i.e., surface density independent of mass for the dark halos (J. Kormendy & K.C. Freeman, in preparation).
4.1.4. Signatures of the internal distribution of
specific angular momentum
The internal distribution of specific angular momentum
(h) of the baryons
(i.e., the mass with specific angular momentum < h) largely
determines the shape of the surface brightness distribution of the disk
rotating in the potential well of the dark matter. Together with
(h) the total
angular momentum and mass of the baryons determine the scale length and
scale surface density of the disk. Therefore the distribution of total
angular momentum and mass for protodisk galaxies determines the observed
distribution of the scale length and scale surface density for disk
galaxies
(Freeman 1970,
de Jong & Lacey
2000).
Many studies have assumed that
(h)
is conserved through the galaxy formation process, most notably
Fall & Efstathiou
(1980).
It is not yet clear if this assumption is correct. Conservation of the
internal distribution of specific angular momentum
(h)
is a much stronger requirement than the conservation of the total
specific angular momentum J / M. Many processes can cause
the internal angular momentum to be redistributed, while leaving the
J / M ratio unchanged. Examples include the effects of
bars, spiral structure
(Lynden-Bell &
Kalnajs 1972)
and internal viscosity
(Lin & Pringle
1987).
The maximum specific angular momentum hmax of the
baryons may be associated with the truncation of the optical disk
observed at about four scale lengths
(de Grijs et al. 2001,
Pohlen et al. 2000).
This needs more investigation. The truncation of disks could be an
important signature of the angular momentum properties of the early
protocloud, but it may have more to do with the critical density for
star formation or the dynamical evolution of the disk. Similarly, in
galaxies with very extended HI, the edge of the HI distribution may give
some measure of hmax in the protocloud. On the other
hand, it may be that the outer HI was accreted subsequent to the
formation of the stellar disk
(van der Kruit 2001), or that the HI edge may just represent the
transition to an ionized disk
(Maloney 1993).
This last item emphasizes the importance of understanding what is going
on in the outer disk. The outer disk offers some potentially important
diagnostics of the properties of the protogalaxy. At present there are
too many uncertainties about the significance of (a) the various
cutoffs in the light and HI distributions, (b) the age gradient
seen by
Bell & de Jong (2000)
from integrated light of disks but not by
Friel (1995)
for open clusters in the disk of the Galaxy, and (c) the
outermost disk being maybe younger but not “zero age,”
which means that there is no real evidence that the disk is continuing
to grow radially. It is possible that the edge of the disk has something
to do with angular momentum of baryons in the protocloud or with disk
formation process, so it may be a useful zero-order or first-order
signature.
4.1.5. Signatures of the CDM hierarchy
CDM predicts a high level of substructure that is in apparent conflict
with observation. Within galaxies, the early N-body simulations appeared
to show that substructure with characteristic velocities in the range 10
< Vc < 30 km s-1 would be destroyed
by merging and virialization of low mass structures
(Peebles 1970,
White 1976,
White & Rees 1978).
It turned out that the lack of substructure was an artefact of the
inadequate spatial and mass resolution
(Moore et al. 1996).
Current simulations reveal 500 or more low mass structures within 300
kpc of an L* galaxy's sphere of influence
(Moore et al. 1999,
Klypin et al. 1999).
This is an order of magnitude larger than the number of low mass
satellites in the Local Group.
Mateo (1998)
catalogues about 40 such objects and suggests that, at most, we are
missing a further 15−20 satellites at low galactic latitude.
Kauffmann et al. (1993)
were the first to point out the satellite problem and suggested that the
efficiency of dynamical friction might be higher than usually
quoted. However, without recourse to fine tuning, this would remove
essentially all of the observed satellites in the Local Group today.
Since the emergence of the CDM paradigm, an inevitable question is
whether a basic building block can be recognized in the near
field. Moore et al. emphasize the self-similar nature of CDM
sub-clustering and point to the evidence provided by the mass spectrum
of objects in rich clusters, independent of the N-body simulations. The
lure of finding a primordial building block in the near field has
prompted a number of tests. If the dark mini-halos comprise discrete
sources, it should be possible to detect microlensing towards a
background galaxy (see “The Dark Halo” above).
The satellite problem appears to be a fundamental prediction of CDM in
the nonlinear regime. Alternative cosmologies have been suggested
involving the reduction of small-scale power in the initial mass power
spectrum
(Kamionkowski &
Liddle 2000),
warm dark matter
(Hogan & Dalcanton
2001,
White & Croft 2000,
Colin et al. 2000),
or strongly self-interacting dark matter
(Spergel &
Steinhardt 2000).
Several authors have pointed out that some of the direct dark-matter
detection experiments are sensitive to the details of the dark matter in
the solar neighborhood.
Helmi et al. (2001)
estimate that there may be several hundred kinematically cold dark
streams passing through the solar neighborhood.
If CDM is correct in detail, then we have simply failed to detect or to
recognize the many hundreds of missing objects throughout the Local
Group. For example, the satellites may be dark simply because baryons
were removed long ago through supernova-driven winds
(Dekel & Silk 1986,
Mac Low & Ferrara
1999).
In support of this idea, X-ray halos of groups and clusters are almost
always substantially enriched in metals ([Fe/H] ≥ -0.5;
Renzini 2000,
Mushotzky 1999).
In fact, we note that up to 70% of the mass fraction in metals is likely
to reside in the hot intracluster and intragroup medium
(Renzini 2000).
Another explanation may be that the absence of baryons in hundreds of
dark satellites was set in place long ago during the reionization
epoch. Many authors note that the accretion of gas on to low-mass halos
and subsequent star formation is heavily suppressed in the presence of a
strong photoionizing background
(Ikeuchi 1986,
Rees 1986,
Babul & Rees 1992,
Bullock et al. 2000).
This effect appears to have a cut-off at low galactic mass at a
characteristic circular velocity close to 30 km s-1
(Thoul & Weinberg
1996,
Quinn et al. 1996),
such that the small number of visible Local Group dwarfs are those that
exceed this cutoff or acquired most of their neutral hydrogen before the
reionization epoch.
Blitz et al. (1999)
suggested that the high-velocity HI gas cloud (HVC) population is
associated with dark mini-halos on megaparsec scales within the Local
Group. This model was refined by
Braun & Burton (1999)
to include only the compact HVCs. The HVCs have long been the subject of
wide-ranging speculation.
Oort (1966)
realized that distances derived from the virial theorem and the HI flux
would place many clouds at Mpc distances if they are
self-gravitating. If the clouds lie at about a Mpc and are associated
with dark matter clumps, then they could represent the primordial
building blocks. However, Hα distances
(Bland-Hawthorn et
al. 1998)
suggest that most HVCs lie within 50 kpc and are unlikely to be
associated with dark matter halos
(Bland-Hawthorn &
Maloney 2001,
Weiner et al. 2001).
We note that several teams have searched for but failed to detect a
faint stellar population in HVCs.
Moore et al.
(1999;
see also
Bland-Hawthorn &
Freeman 2000)
suggested that ultrathin disks in spirals are a challenge to the CDM
picture in that disks are easily heated by orbiting masses. However,
Font et al. (2001)
found that in their CDM simulations, very few of the CDM sub-halos come
close to the optical disk.
At present there are real problems in reconciling the predictions of CDM
simulations with observations on scales of the Local Group.
4.2. First Order Signatures – Information
Preserved Since the Main Epoch of Baryon Dissipation
4.2.1. The structure of the disk
At what stage in the evolution of the disk are its global properties
defined? In part, we have already discussed this question in
“Signatures of the Internal Distribution of Specific Angular
Momentum” above. The answer depends on how the internal angular
momentum distribution
(h)
has evolved as the disk dissipated and various nonaxisymmetric features
like bars and spiral structure came and went. Viscous processes
associated with star formation, as suggested by
Lin & Pringle
(1987),
may also contribute to the evolution of the
(h) distribution.
The global structure of disks is defined by the central surface
brightnes Io and the radial scalelength h of
the disk.
de Jong & Lacey
(2000)
evaluated the present distribution of galaxies in the
(Io, h) plane (Figure 4). If
(h)
has indeed remained roughly constant, as is often assumed for
discussions of disk formation (e.g.,
Fall & Efstathiou
1980,
Fall 1983),
then the global parameters of the disk – the scale length,
central surface brightness and the Tully-Fisher relation – are
relics of the main epoch of baryon dissipation.
Figure 4. The density distribution of Sa to
Sm galaxies over effective radius re and effective
surface brightness µe. The top panel shows the raw
(unweighted) distribution and the bottom panel shows the luminosity
weighted distribution
(de Jong & Lacey
2000).
4.2.2. Can disks preserve fossil information?
Here, we consider radial and vertical fossil gradients in the disk, in
particular of abundance and age. Our expectation is that much of the
information will be diluted through the dynamical evolution and radial
mixing of the disk.
For spirals, different mechanisms may be at work to establish gradients
(Molla et al. 1996):
(a) a radial variation of the yield due either to the stellar
metal production or to the initial mass function, (b) a radial
variation of the timescale for star formation, (c) a radial
variation for the timescale of infall of gas from outside the disk. Once
gradients are established, these can be amplified or washed out by
radial mixing
(Edmunds 1990,
Goetz & Koeppen
1992).
Most stars are born in large clusters numbering hundreds or even
thousands of stars. Some clusters stay together for billions of years,
whereas others become unbound shortly after the initial starburst,
depending on the star formation efficiency. When a cloud disperses, each
star suffers a random kick superimposed on the cloud's mean
motion. Thereafter, stars are scattered by transient spiral arm
perturbations and star-cloud encounters.
These perturbations allow the star to migrate in integral space. During
interaction with a single spiral event of pattern speed
Ωp, a star's energy and angular momentum change while
it conserves its Jacobi integral: In the (E, J) plane,
stars move along lines of constant IJ = E -
Ωp J. The star undertakes a random walk in the
(E, J) plane, perturbed by a series of spiral arm events
(Sellwood 1999,
Dehnen 2000).
N-body models of disk evolution indicate that radial mixing is strong
(Sellwood 2001,
Lynden-Bell &
Kalnajs 1972).
This is believed to be driven by transient spiral waves that heat
the in-plane motions, although the process is not yet well
understood. Long-term spiral arms produce no net effect. Remarkably, a
single spiral wave near co-rotation can perturb the angular
momentum of a star by ∼ 20% without significant heating: The star is
simply moved from one circular orbit to another, inwards or outwards, by
up to 2 kpc
(Sellwood & Kosowsky 2000).
Substantial variations in the angular momentum of a star are possible
over its lifetime.
In addition to radial heating, stars experience vertical disk heating:
Their vertical velocity dispersion increases as they age. This is
believed to occur through a combination of in-plane spiral-arm heating
and scattering off giant molecular clouds (e.g.,
Spitzer &
Schwarzschild 1953,
Carlberg & Sellwood
1985).
The in-plane heating is most effective at the inner and outer Lindblad
resonances and vanishes at corotation. In the vertical direction, an
age-velocity dispersion relation is observed for stars younger than
about 3 Ga, but older disk stars show a velocity dispersion that is
independent of age (Figure 5). Thus, the
vertical structure does depend on the mean age of the population for
τ < 3 Ga
(Edvardsson et al. 1993,
confirmed from Hipparcos data by
Gomez et al. 1997).
Figure 5. The relation between the three
components of the velocity dispersion and the stellar age, as derived
by
Quillen & Garnett
(2001)
for stars from the sample of
Edvardsson et
al. (1993).
Stars with ages between 2 and 10 Ga belong to the old thin disk: Their
velocity dispersion is independent of age. The younger stars show a
smaller velocity dispersion. The velocity dispersion doubles abruptly
at an age of about 10 Ga; these older stars belong to the thick
disk.
As the amplitude of the random motions increases, the star becomes less
vulnerable to heating by transient spiral waves, and the heating process
is expected to saturate. This probably happens after about 3 Ga
(Binney & Lacey 1988,
Jenkins & Binney
1990),
consistent with observation. This is important for our purpose here. It
means that dynamical information is preserved about the state of the
thin disk at an early epoch, or roughly τL - 3
≈ 7 Ga ago, for which τL is the look-back time
when the disk first began to form.
The survival of old open clusters like NGC 6791, Berkeley 21 and
Berkeley 17
(Friel 1995,
van den Bergh 2000)
is of interest here. The oldest open clusters exceed 10 Ga in age and
constitute important fossils
(Phelps & Janes
1996).
Both old and young open clusters are part of the thin disk. If the
heating perturbations occur over a lengthscale that significantly
exceeds the size of an open cluster, it seems likely that the cluster
will survive. A large spiral-arm heating event will heat many stars
along their IJ trajectories. The trace of the heating event
is likely to survive for a very long time but be visible only in
integral space
(Sellwood 2001).
We note that vertical abundance gradients have not been seen among the
open clusters
(Friel & Janes 1993).
About 4% of disk stars are super metal-rich (SMR) relative to the Hyades
(Castro et al. 1997).
SMR stars of intermediate age appear to have formed a few kpc inside of
the Solar circle from enriched gas. The oldest SMR stars appear to come
from the Galactic Center: Their peculiar kinematics and outward
migration may be associated with the central bar
(Carraro et al. 1998,
Grenon 1999).
In summary, our expectation is that fossil gradients within the disk are
likely to be weak. This is borne out by observations of both the stars
and the gas
(Chiappini et al. 2001).
The vertical structure of the disk preserves another fossil – the
thick disk – which we discuss in the next section. Like the open
clusters, this component also does not show a vertical abundance
gradient
(Gilmore et al. 1995).
In later sections, we argue that this may be the most important fossil
to have survived the early stages of galaxy formation.
4.2.3. Disk heating by accretion: the thick
disk
Heating from discrete accretion events also imposes vertical structure
on the disk
(Quinn & Goodman
1986,
Walker et al. 1996).
Such events can radically alter the structure of the inner disk and the
bulge (see Figure 3d for an
example) and are currently believed to have generated the thick disk of
the Galaxy.
The galactic thick disk was first recognized by
Gilmore & Reid
(1983).
It includes stars with a wide range of metallicity, from -2.2 ≤
[Fe/H] ≤ -0.5
(Chiba & Beers 2000):
Most of the thick disk stars are in the more metal-rich end of this
range. The velocity ellipsoid of the thick disk is observed to be
(σR, σφ, σz =
(46 ± 4, 50 ± 4, 35 ± 3) km s-1 near the sun,
with an asymmetric drift of about 30 km s-1. For comparison,
the nearby halo has a velocity ellipsoid (σR,
σφ, σz = (141 ± 11,
106 ± 9, 94 ± 8) km s-1 and its asymmetric drift is
about 200 km s-1.
The mean age of the thick disk is not known. From photometric age-dating
of individual stars, the thick disk appears to be as old as the globular
clusters. Indeed, the globular cluster 47 Tuc (age 12.5 ± 1.5 Ga;
Liu & Chaboyer 2000)
is often associated with the thick disk.
After
Quinn & Goodman
(1986),
Walker et al. (1996)
showed in detail that a low mass satellite could substantially heat the
disk as it sinks rapidly within the potential well of a galaxy with a
live halo. The conversion of satellite orbital energy to disk thermal
energy is achieved through resonant scattering. Simulations of satellite
accretion are important for understanding the survival of the thin disk
and the origin of the thick disk. This is particularly relevant within
the context of CDM. The satellites which do the damage are those that
are dense enough to survive tidal disruption by the Galaxy. We note that
even dwarf spheroidals which appear fluffy are in fact rather dense
objects dominated by their dark matter (J. Kormendy & K. Freeman, in
preparation).
It is fortuitous that the Galaxy has a thick disk, since this is not a
generic phenomenon. The disk structure may be vertically stepped as a
consequence of past discrete accretion events. The
Edvardsson et al. (1993)
data (Figure 5) appears to show an abrupt
increase in the vertical component of the stellar velocity dispersion at
an age of 10 Ga; see also
Strömgren (1987).
Freeman (1991)
argued that the age−velocity dispersion relation shows three
regimes: stars younger than 3 Ga with σz ∼ 10 km
s-1, stars between 3 and 10 Ga with σz ∼
20 km s-1, and stars older than 10 Ga with
σz ∼ 40 km s-1. The first regime
probably arises from the disk heating process due to transient spiral
arms which we described in the previous section. The last regime is the
thick disk, presumably excited by an ancient discrete event.
Can we still identify the disrupting event that led to the thick disk?
There is increasing evidence now that the globular cluster ω Cen
is the stripped core of a dwarf elliptical (see “Globular
Clusters” below). It is possible that the associated accretion
event or an event like it was the event that triggered the thick disk to
form.
In summary, it seems likely that the thick disk may provide a
snap-frozen view of conditions in the disk shortly after the main epoch
of dissipation. Any low level chemical or age gradients would be of
great interest in the context of dissipation models. In this regard,
Hartkopf & Yoss
(1982)
argued for the presence of a vertical abundance gradient in the thick
disk, although Gilmore, Wyse & Jones (1995) found no such
effect. Because stars of the thick disk spend relatively little time
near the galactic plane, where the spiral arm heating and scattering by
giant molecular clouds is most vigorous, radial mixing within the thick
disk is unlikely to remove all vestiges of a gradient. If our earlier
suggestions are right (see “Signatures of the Internal
Distribution of Specific Angular Momentum” above), we might
expect to see a different truncation radius for the thick disk compared
to the thin disk.
4.2.4. Is there an age-metallicity relation?
Some fossil information has likely been preserved since the main epoch
of baryon dissipation. The inner stellar bulge is a striking example. It
is characterized by old, metal-rich stars, which seems to be at odds
with the classical picture where metals accumulate with time
(Tinsley 1980).
However, the dynamical timescales in the inner bulge are very short
compared to the outer disk and would have allowed for rapid enrichment
at early times. This is consistent with the frequent occurrence of
metal-rich cores of galaxies observed at high redshift
(Hamann & Ferland
1999).
The dynamical complexity of the Galactic bulge may not allow us to
determine the sequence of events that gave rise to it. We anticipate
that this will come about from far-field cosmology
(Ellis et al. 2000).
The existence of an age-metallicity relation (AMR) in stars is a very
important issue, about which there has long been disagreement.
Twarog (1980),
Meusinger et al. (1991)
provide evidence for the presence of an AMR, while Carlberg et
al. (1985) find that the metallicity of nearby F stars is approximately
constant for stars older than about 4 Ga. More recently it has become
clear that an AMR is apparent only in the solar neighborhood and is
strictly true only for stars younger than 2 Ga and hotter than log
Teff = 3.8
(Feltzing et al. 2001).
Edvardsson et al. (1993)
demonstrate that there is no such relation for field stars in the old
disk. Similarly,
Friel (1995)
shows that there is no AMR for open clusters (see “Open
Clusters” below): she goes on to note that
Apparently, over the entire age of the disk, at any position in the
disk, the oldest clusters form with compositions as enriched as those of
much younger objects.
In fact, it has been recognized for a long time (e.g.,
Arp 1962,
Eggen & Sandage 1969,
Hirshfeld et al. 1978)
that old, metal-rich stars permeate the galaxy, throughout the disk, the
bulge and the halo. We regard the presence of old metal-rich stars as a
first-order signature. An age-metallicity relation which applies to all
stars would have been an important second-order signature, but we see no
evidence for such a relation, except among the young stars.
4.2.5. Effects of environment and internal
evolution
Environmental influences are operating on all scales of the hierarchy
and across all stages of our signature classification, so our attempts
to classify signatures are partly artificial. Within CDM, environmental
effects persist throughout the life of the galaxy.
The parameters that govern the evolution of galaxies are among the key
unknowns of modern astrophysics. Are the dominant influences internal
(e.g., depth of potential) or external (e.g., environment) to galaxies?
We consider here the effects of environment and internal evolution on
the validity of the first-order signatures of galaxy formation (i.e.,
the properties that may have been conserved since the main epoch of
baryon dissipation).
The well-known G dwarf problem indicates that external influences are
important. A simple closed box model of chemical evolution predicts far
too many metal-poor stars in the solar neighborhood
(Tinsley 1980).
This problem is easily remedied by allowing gas to flow into the region
(Lacey & Fall 1983,
1985,
Clayton 1987,
1988,
Wyse & Silk 1989,
Matteuci & Francois
1989,
Worthey et al. 1996).
In the context of CDM, this is believed to arise from the continued
accretion of gas-rich dwarfs (e.g.,
Cole et al. 1994,
Kauffmann & Charlot
1998).
Environment is clearly a key factor. Early type galaxies are highly
clustered compared to late type galaxies
(Hubble & Humason
1931,
Dressler 1980).
Trager et al. (2000)
find that for a sample of early-type galaxies in low-density
environments, there is a large spread in the Hβ index (i.e., age),
but little variation in metallicity. For galaxies in the Fornax cluster,
Kuntschner (2000)
finds the opposite effect: A large spread in metallicity is present with
little variation in age. This probably reflects strong differences in
environment between the field and the cluster.
Another likely environmental effect is the fraction of S0 galaxies in
clusters, which shows a rising trend with redshift since z
≈ 0.4
(Jones et al. 2000).
Furthermore, S0 galaxies in the Ursa Major cluster show age gradients
that are inverted compared to field spirals, in the sense that the cores
are young and metal-rich
(Tully et al. 1996,
Kuntschner & Davies
1998).
Both of these effects involve more recent phenomena and would be
properly classified as second-order signatures.
Internal influences are also at work. A manifestation is the
color-magnitude relation (CMR) in early-type
(Sandage &
Visvanathan 1978)
and late-type
(Peletier & de Grijs
1998)
galaxies. The CMR does not arise from dust effects
(Bell & de Jong 2000)
and must reflect systematic variations in age and/or metallicity with
luminosity. In the case of ellipticals, the CMR is believed to reflect a
mass-metallicity dependence
(Faber 1973,
Bower et al. 1998).
The relation is naturally explained by supernova-driven wind models in
which more massive galaxies retain supernova ejecta and thus become more
metal rich and redder
(Larson 1974,
Arimoto & Yoshii
1987).
The CMR is presumably established during the main phase of baryon
dissipation and is a genuine first-order signature.
Concannon et al. (2000)
analyzed a sample of 100 early-type galaxies over a large range in
mass. They found that lower-mass galaxies exhibit a larger range in age
than higher-mass galaxies. This appears to show that smaller galaxies
have had a more varied star formation history, which is at odds with the
naive CDM picture of low-mass galaxies being older than high-mass
galaxies
(Baugh et al. 1996,
Kauffman 1996).
The work of
Concannon et al. (2000)
shows the presence of a real cosmic scatter in the star formation
history. It is tempting to suggest that this cosmic scatter relates to
different stages of evolution within the hierarchy. In this sense, we
would regard the Concannon et al. result as a first-order manifestation
of galaxy formation (see “Timescales and Fossils” above).
Spiral galaxies commonly show color gradients that presumably reflect
gradients in age and metallicity
(Peletier & de Grijs
1998).
Faint spiral galaxies have younger ages and lower metallicities relative
to bright spirals. In a study of 120 low-inclination spirals,
Bell & de Jong (2000)
found that the local surface density within galaxies is the most
important parameter in shaping their star formation and chemical
history. However, they find that metal-rich galaxies occur over the full
range of surface density. This fact has a remarkable resonance with the
distribution of the metal-rich open clusters that are found at any
position in the Galactic disk (see “Is There an Age-Metallicity
Relation?” above). Bell & de Jong argue that the total mass
is a secondary factor that modulates the star formation history. Once
again, these authors demonstrate the existence of cosmic scatter that
may well arise from variations in environment.
4.3. Second Order Signatures – Major Processes
Involved in Subsequent Evolution
Here we consider relics of processes that have taken place in the Galaxy
since most of the baryonic mass settled to the disk. There are several
manifestations of these processes, probably the most significant of
which is the star formation history of the disk, for which the open
clusters are particularly important probes.
There is a wealth of detail relating to anomalous populations throughout
the Galaxy, discussed at length by
Majewski (1993).
Examples include an excess of stars on extreme retrograde orbits
(Norris & Ryan 1989,
Carney et al. 1996),
metal-poor halo stars of intermediate age
(Preston et al. 1994)
and metal-rich halo A stars
(Rodgers et al. 1981).
In an earlier section, we discussed observational signatures of the CDM
hierarchy in the Galactic context. In fact, detailed observations in
velocity space are proving to be particularly useful in identifying
structures that have long since dispersed in configuration space. In
external galaxies, related structures are showing up as low surface
brightness features. We do not know what role globular clusters play in
the galaxy formation picture, but we include them here because at least
one of them appears now to be the nucleus of a disrupted dwarf galaxy.
The star formation history (SFH) of our Galaxy has been very difficult
to unravel. Derived star formation histories range from a roughly
uniform star formation rate over the history of the disk to a SFH that
was highly peaked at early times (e.g.,
Twarog 1980,
Rocha-Pinto et al. 2000,
Just 2001).
Galaxies of the Local Group show a great diversity in SFH
(Grebel 2001),
although the average history over the Local Group appears consistent
with the mean cosmic history
(Hopkins et al. 2001).
The present emphasis is on star formation studies that make use of the
integrated properties of external galaxies, but it should be noted that
this is necessarily weighted towards the most luminous populations. Key
results for external galaxies are reviewed in “Effects of
Environment and Internal Evolution” above. It was concluded that
environmental effects are very significant in determining the SFH for
individual galaxies.
The conventional approach to the study of chemical evolution in galaxy
disks is to consider the solar neighborhood a closed box, and to assume
that it is representative of all disks. Simple mathematical formulations
have developed over the past 40 years
(van den Bergh 1962,
Schmidt 1963,
Pagel & Patchett
1975,
Talbot & Arnett
1971,
Tinsley 1980,
Twarog 1980,
Pitts & Tayler
1989).
Most observations are interpreted within this framework. The SFH is
quantified in terms of stellar age, stellar (+gas) metallicity and, to a
lesser extent, the existing gas fraction.
The use of broadband photometry coupled with stellar population
synthesis is a well-established technique for probing the SFH of galaxy
populations from integrated light. The power of the method is its
simplicity, although it cannot uniquely disentangle the age-metallicity
degeneracy
(Bica et al. 1990,
Charlot & Silk 1994).
Another widely used technique is the Lick index system
(Burstein et al. 1984)
further refined in
Worthey et al. (1994),
Trager et al. (1998).
In this system, the Hβ index is the primary age-sensitive spectral
indicator, whereas the Mg and Fe indices are the primary metallicity
indicators. The Lick indices have well-known limitations: They
correspond to low spectroscopic resolution (8−9 Å), require
difficult corrections for internal galaxy motions, and are not
calibrated onto a photometric scale. Furthermore, two of the most
prominent Lick indices – Mg2 λ5176 and Fe
λ5270 – are now known to be susceptible to contamination
from other elements, in particular Ca and C
(Tripicco & Bell
1995).
How best to measure galaxy ages is a subject with a long history. The
most reliable methods to date involve the low order transitions
(n < 4) of the Balmer series. Ages derived from the Hγ
equivalent width have been used by
Jones & Worthey
(1995).
Rose (1994),
Caldwell & Rose
(1998)
have pioneered the use of even higher-order Balmer lines to break the
age-metallicity degeneracy
(Worthey 1994).
These higher-order lines are less affected by Balmer line emission from
the interstellar medium. They develop a line ratio index Hn/Fe
which is a sum over Hγ, Hδ and H8 lines with respect to
local Fe lines. The most recent demonstration of the power of this index
can be found in
Concannon et al. (2000).
Ultimately, full spectrum fitting matched to spectral synthesis models
holds the most promise
(Vazdekis 1999).
The new models, which have a fourfold increase in spectroscopic
resolution compared to the Lick system, show that the isochrone or
isochemical grid lines overlaid on a plot of two Lick indices are more
orthogonal than the Worthey models. Thus, galaxies like NGC 4365 that
exhibit no age gradient in the Vazdekis models
(Davies et al. 2001;
see Figure 6a) appear to show an age
spread in the Worthey models. Interestingly, NGC 4150 exhibits an
abundance spread with constant age
(Figure 6b).
Figure 6. Sauron integral field
observations of NGC 4365 (top panels) and NGC 4150 (bottom
panels). Left panels: reconstructed surface brightness maps. Right
panels: Hβ versus [MgFe5270] diagram. The points were derived
from the Sauron datacubes by averaging along the corresponding
color-coded isophotes (Bacon et al. 2001,
Davies et al. 2001).
The [Fe/H] vs. age grid is derived from
Vazdekis (1999).
[We acknowledge Harald
Kuntschner and the Sauron team for these images.]
4.3.3. Low surface brightness structures in
galaxies
Dynamical interaction between galaxies led to a range of structures
including stellar shells
(Malin & Carter
1980,
Quinn 1984),
fans
(Weil et al. 1997),
and tidal streamers
(Gregg & West 1998,
Calcaneo-Roldan et
al. 2000,
Zheng et al. 1999).
Some excellent examples are shown in
Figure 7. We see evidence of multiple nuclei,
counter-rotating cores, and gas in polar orbits. At low light levels,
the outermost stellar contours of spiral disks appear frequently to
exhibit departures from axisymmetry
(Rix & Zaritsky
1995).
The same is true for spiral arms in all Hubble types
(Schoenmakers et
al. 1997;
Cianci 2002).
Figure 7. Examples of normal spirals with
faint stellar streamers in the outer halo (see text): (a) M104
where the streamer is on a much larger scale than shown in
Figure 3(c)
(Malin & Hadley
1997);
(b) M83
(Malin & Hadley
1997);
(c) NGC 5907
(Shang et al. 1998);
(d) M31
(Ibata et
al. 2001a).
The stellar streamers are particularly interesting, as these may provide
important constraints on galaxy models, particularly as kinematic
measurements become possible through the detection of planetary
nebulae. More than a dozen stellar streams are already known and this is
probably indicative of a much larger population at very low surface
brightness.
Johnston et al. (2001)
show that stellar streamers can survive for several gigayears and are
only visible above the present optical detection limit
(µV = 30 mag arcsec-2) for roughly 4 ×
108 yr. A few galaxy groups (e.g., the Leo group) do show
large-scale HI filaments that can remain visible for many Ga.
Deep CCD imaging has revealed a stellar loop around NGC 5907
(Shang et al. 1998)
and a stellar feature extending from NGC 5548
(Tyson et al. 1998).
The technique of photographic amplification has revealed stellar
streamers in about ten sources
(Malin & Hadley
1997,
Calcaneo-Roldan et
al. 2000,
Weil et al. 1997).
For these particular observations, the limiting surface brightness is
µV ≈ 28.5 mag arcsec-2. For all
of these systems, we estimate that the total stream luminosities are in
the range 3−20 × 107 L⊙.
In a recent development, wide-field CCD cameras have revealed stellar
streamers through multiband photometry of millions of individual
sources. A pointillist image can then be reconstructed in narrow color
intervals so as to enhance features with respect to the field. This has
led to the discovery of a stellar stream in M31
(Ibata et al. 2001a)
and tidal tails extending from the globular cluster Pal 5
(Odenkirchen et
al. 2001).
This technique has the potential to push much deeper than the direct
imaging method described above.
The low surface brightness universe is notoriously difficult to
observe. Modern telescope and instrument designs are simply not
optimized for this part of parameter space. Many claims of diffuse light
detections in the neighborhood of galaxies have been shown to arise from
scattered light internal to the instrument.
In “Structures in Phase Space,” below, we discuss moving
groups identified within the Galaxy from proper motion and spectroscopic
surveys. Their projected surface brightness is µV =
30−34 mag arcsec-2, below the limit of modern imaging
techniques.
Looking farther afield, we see evidence for discrete accretion events in
the making. The Galaxy is encircled by satellite galaxies that appear
confined to one or two great streams across the sky
(Lynden-Bell &
Lynden-Bell 1995).
The most renowned of these are the Magellanic Clouds and the associated
HI Magellanic stream. All of these are expected to merge with the Galaxy
in the distant future, largely due to the dynamical friction from the
extended halo.
In the context of near-field cosmology, we believe that the thick disk
and the old open clusters of the thin disk are among the most important
diagnostics. The open clusters are the subject of an outstanding and
comprehensive review by
Friel (1995).
Here, we summarize the properties that are most important for our purpose.
Both old and young clusters are part of the thin disk. Their key
attribute is that they provide a direct time line for investigating
change, which we explore in “The Gaiasphere and the Limits of
Knowledge,” below. The oldest open clusters exceed 10 Ga in age
and constitute important fossils
(Phelps & Janes
1996).
In “Can disks preserve fossil information,” we noted that
the survival of these fossil clusters is an interesting issue in its own
right.
Friel (1995)
finds no old open clusters within a galactocentric radius of 7 kpc;
these are likely to have disrupted or migrated out of the central
regions
(van den Bergh &
McClure 1980).
It has long been recognized that open clusters walk a knife edge between
survival and disruption
(King 1958a,
b,
c).
Like field stars in the disk,
Janes & Phelps
(1994)
find that the old cluster population (relative to Hyades) is defined by
a 375-pc scale height exponential distribution, whereas young clusters
have a 55-pc scale height
(Figure 8a,b). Again, like the
field stars, vertical abundance gradients have not been seen in open
clusters
(Friel & Janes 1993),
although radial gradients are well established
(Friel 1995,
van den Bergh 2000).
For old open clusters,
Twarog et al. (1997)
claim evidence for a stepped radial metallicity distribution where
[Fe/H] ≈ 0 within 10 kpc, falling to [Fe/H] ≈ -0.3 in the
outer disk. However, this effect is not seen in young objects, e.g., HII
regions and B stars
(Henry 1998).
Figure 8. (a) The distribution of
open clusters younger than Hyades with height from the plane as a
function of Galactocentric distance Rgc
(Friel 1995).
The Sun is at 8.5 kpc. (b) The distribution of clusters with
ages equal to or greater than the Hyades. (c) The open clusters
exhibit a well-defined abundance gradient. (d) There is no
discernible age-metallicity relation (AMR) when the cluster abundances
are corrected for the radial abundance gradient.
In Figure 8c, both the old and young open
clusters show essentially the same radial trend in metallicity. After
reviewing the available observations,
Friel (1995)
finds no evidence for an age-metallicity relation for open clusters
(Figure 8d). In agreement with
Eggen & Sandage
(1969),
she notes that over the entire age of the disk, at any position in the
disk, the oldest clusters form with compositions as enriched as those of
much younger objects.
These remarkable observations appear to indicate that shortly after the
main epoch of baryon dissipation, the thin disk was established at least
as far out as 15 kpc. The oldest open clusters approach the age of the
thick disk. Since, in “Disk Heating by Accretion,” we
noted that the thick disk is likely to be a snap frozen picture of the
thin disk shortly after disk formation, we would expect the truncation
of the thick disk (see “Signatures of Global Quantities,”
above) to reflect the extent of the thin disk at the epoch of the event
that puffed up the thick disk.
We have long suspected that globular clusters are the fossil remnants of
violent processes in the protogalactic era
(Peebles & Dicke
1968).
But there is a growing suspicion that globulars are telling us more
about globulars than galactic origins
(Harris 2001).
The Milky Way has about 150 globular clusters with 20% lying within a
few kiloparsecs of the Galactic Center. They constitute a negligible
fraction of the light and mass (2%) of the stellar halo today. Their
significance rests in their age. The oldest globular clusters in the
outer halo have an age of 13 ± 2.5 Ga (90% confidence).
The ages of the oldest globular clusters in the inner and outer halo,
the Large Magellanic Cloud and the nearby
Fornax and Sgr dwarf spheroidal galaxies show a remarkable
uniformity. To a precision of
± 1 Ga, the onset of globular cluster formation was well
synchronized over a volume centered on our Galaxy with a radius > 100
kpc
(Da Costa 1999).
Globular cluster stars are older than the oldest disk stars, e.g., white
dwarfs and the oldest red giants. These clusters are also more metal
poor than the underlying halo light in all galaxies and at all radii
(Harris 1991),
but again there are exceptions to the rule. Since Morgan's (1950) and
Kinman's (1959)
classic work, we have known that there are two distinct populations of
globular clusters in the Galaxy. The properties that we associate with
these two populations today were derived by
Zinn (1985)
who showed that they have very different structure, kinematics and
metallicities. The halo population is metal poor ([Fe/H] < -0.8) and
slowly rotating with a roughly spherical distribution; the disk
population is metal rich ([Fe/H] > -0.8) and in rapid rotation.
A major development has been the discovery of young globular clusters in
disturbed or interacting galaxies, e.g., NGC 1275
(Holtzman et al. 1992),
NGC 7252
(Whitmore et al. 1993)
and the Antennae
(Whitmore &
Schweizer 1995).
Schweizer (1987)
first suspected that globular clusters were formed in mergers. Later,
Ashman & Zepf (1992)
predicted that the HST would reveal young globular clusters through
their compact sizes, high luminosities and blue colors. The very high
internal densities of globular clusters today must partly reflect the
conditions when they were formed.
Harris & Pudritz
(1994)
present a model for globular clusters produced in fragmenting giant
molecular clouds, which are of the right mass and density range to
resemble accretion fragments in the Searle-Zinn model.
Globular clusters have been elegantly referred to as “canaries in
a coal mine”
(Arras & Wasserman
1999).
They are subject to a range of disruptive effects, including two-body
relaxation and erosion by the tidal field of their host galaxy, and the
tidal shocking that they experience as their orbits take them through
the galactic disk and substructure in the dark halo. In addition to
self-destruction through stellar mass loss, tidal shocking may have been
very important in the early universe
(Gnedin et al. 1999).
If globular clusters originally formed in great numbers, the disrupted
clusters may now contribute to the stellar halo
(Norris & Ryan 1989,
Oort 1965).
Halo field stars and globular clusters in the Milky Way have similar
mean metallicities
(Carney 1993);
however, the metallicity distribution of the halo field stars extends to
much lower metallicity ([Fe/H] ≃ -5) than that of the globular
clusters ([Fe/H] ≃ -2.2). We note again the remarkable similarity
in the metallicity range of the globular clusters and the thick disk
(-2.2 ≲ [Fe/H] ≲ -0.5).
In the nucleated dwarf elliptical galaxies
(Binggeli et al. 1985),
the nucleus typically provides about 1% of the total luminosity;
globular clusters could be considered as the stripped nuclei of these
satellite objects without exceeding the visible halo mass
(Zinnecker & Cannon
1986,
Freeman 1993).
It is an intriguing prospect that the existing globular clusters could
be the stripped relicts of an ancient swarm of protogalactic stellar
fragments, i.e., the original building blocks of the Universe.
In the Searle-Zinn picture, globular clusters are intimately linked to
gas-rich, protogalactic infalling fragments. Multiple stellar
populations have recently been detected in ω Cen, the most massive
cluster in the Galaxy
(Lee et al. 1999).
How did ω Cen retain its gas for a later burst? It now appears
that it was associated with a gas-rich dwarf, either as an in situ
cluster or as the stellar nucleus. The present-day cluster density is
sufficiently high that it would have survived tidal disruption by the
Galaxy, unlike the more diffuse envelope of this dwarf galaxy. The very
bound retrograde orbit supports the view that ω Cen entered the
Galaxy as part of a more massive system whose orbit decayed through
dynamical friction.
If globular clusters are so ancient, why are the abundances of the most
metal-poor population as high as they are? Because it does not take much
star formation to increase the metal abundance up to [Fe/H] = -1.5
(Frayer & Brown
1997),
the cluster abundances may reflect low levels of star formation even
before the first (dark + baryon) systems came together.
Old age is not necessarily associated with low metallicity (compare
“Timescales and Fossils” above). We recall that CO has
been detected at z ∼ 5
(Yun et al. 2000).
Hamann & Ferland
(1999)
demonstrate that stellar populations at the highest redshift currently
observed appear to have solar or super-solar metallicity. We believe
that there is no mystery about high abundances at high redshift. The
dynamical times in the cores of these systems are short, so there has
been time for multiple generations of star formation and chemical
enrichment. In this sense, the cores of high redshift galaxies need not
be relevant to the chemical properties of the globular clusters,
although both kinds of objects were probably formed at about the same
time.
The first generation of globular clusters may have been produced in
merger-driven starbursts when the primordial fragments came together for
the first time. If at least some fragments retained some of their
identity while the halo was formed, a small number of enrichment events
per fragment would ensure a Poissonian scatter in properties between
globular clusters, and multiple populations within individual clusters
(Searle & Zinn
1978).
4.3.6. Structures in phase space
One class of systems that exhibit coherence in velocity space are the
open clusters associated with the disk. Here the common space motion of
the stars with respect to the Sun is perceived as a convergence of the
proper motions to a single point (strictly speaking, minimum volume) on
the sky
(Boss 1908;
see
de Zeeuw et al. 1999
for a recent application). More than a dozen such systems have been
identified this way. However, these are all young open clusters largely
associated with the Gould belt. With sufficiently precise kinematics, it
may be possible to identify open clusters that have recently dispersed,
particularly if the group is confined to a specific radial zone by
resonances in the outer disk. For example,
Feltzing & Holmberg
(2000)
show that the metal-rich ([Fe/H] ≈ 0.2) moving group HR 1614,
thought to be 2 Ga old, can be identified in the Hipparcos data set.
Recently, attention has turned to a diverse set of moving groups that
are thought to be associated with the stellar halo and in some instances
are clearly fossils associated with accretion events in the distant
past. The evidence for these groups dates back to the discovery of the
halo itself. Shortly before the publication of the landmark ELS paper,
Eggen & Sandage
(1959)
discovered that the nearby high-velocity star, Groombridge 1830, belongs
to a moving group now passing through the Galactic disk.
In a long series of papers, Eggen went on to identify a number of moving
groups, some of which appear to encompass the solar neighborhood, and
others that may be associated with the halo. The relevant references are
given by
Taylor (2000).
Various authors have noted that many of the groups are difficult to
confirm
(Griffin 1998,
Taylor 2000).
More systematic surveys over the past few decades have identified a
number of moving populations associated with the halo
(Freeman 1987,
Majewksi 1993),
although the reality of some of these groups is still debated. The
reality of these groups is of paramount importance in the context of
halo formation.
Majewski et al. (1996)
suspect that much or all of the halo could exhibit phase-space clumping
with data of sufficient quality.
In recent years, the existence of kinematic sub-structure in the
galactic halo has become clear.
Helmi et al. (1999)
identified 88 metal-poor stars within 1 kiloparsec of the Sun from the
Hipparcos astrometric catalogue. After deducing accurate 3-D space
motions, they found a highly significant group of 8 stars that appear
clumped in phase space and confined to a highly inclined orbit.
The most dramatic evidence is surely the highly disrupted
Sgr dwarf galaxy identified by
Ibata et al. (1994,
1995).
These authors used multi-object spectroscopy to uncover an elongated
stellar stream moving through the plane on the far side of the
Galaxy. The Sgr dwarf is a low mass dwarf spheroidal galaxy
about 25 kpc
from the Sun that is presently being disrupted by the Galactic tidal
field. The long axis of the prolate body (axis ratios ∼ 3:1:1) is
about 10 kpc, oriented perpendicular to the Galactic plane along
ℓ = 6∘ and centered at b =
-15∘. Sgr contains a mix of stellar
populations, an extended dark halo (mass ≥ 109
M⊙) and at least four globular clusters
(Ibata et al. 1997).
The Sgr stream has since been recovered by several photometric surveys
(Vivas et al. 2001,
Newberg et al. 2002,
Ibata et al. 2001c).
N-body simulations have shown that stellar streams are formed when low
mass systems are accreted by a large galaxy (e.g.,
Harding et al. 2001).
Streamers remain dynamically cold and identifiable as a kinematic
substructure long after they have ceased to be recognizable in star
counts against the vast stellar background of the galaxy
(Tremaine 1993,
Ibata & Lewis 1998,
Johnston 1998,
Helmi & White 1999).
Within the Galaxy, moving groups can be identified with even limited
phase-space information
(de Bruijne 1999,
de Zeeuw et al. 1999).
This also holds for satellites orbiting within the spherical halo, since
the debris remains in the plane of motion for at least a few orbits
(Lynden-Bell &
Lynden-Bell 1995,
Johnston et al. 1996).
But a satellite experiencing the disk potential no longer conserves its
angular momentum and its orbit plane undergoes strong precession
(Helmi & White
1999).
In Figure 9, we show the sky projection of a
satellite 8 Ga after disruption. These more complex structures are
usually highly localized and therefore easy to recognize in the space of
conserved quantities like energy and angular momentum for individual
stars.
The evolution in phase space of a disrupting satellite is well behaved
as its stars become phase mixed. Its phase space flow obeys Liouville's
theorem, i.e., the flow is incompressible. Highly intuitive accounts are
given elsewhere
(Carlberg 1986,
Tremaine 1999,
Hernquist & Quinn
1988).
It should be possible to recognize partially phase-mixed structures that
cover the observed space, although special techniques are needed to find
them.
Four astrometric space missions are planned for the next decade. These
are the proposed German DIVA mission (∼ 2003); the FAME mission
(∼ 2005) and the pointed SIM mission (∼ 2005); and the ESA Gaia
mission (∼ 2009) which will observe a billion stars to V ∼ 20,
with accuracy 10µas at a V ∼ 15. The web sites for these
missions are at:
http://www.ari.uni-heidelberg.de/diva/
The astrometric missions will derive 6-dimensional phase space coordinates and spectrophotometric properties for millions of stars within a 20 kiloparsec sphere – the Gaiasphere. The ambitious Gaia mission will obtain distances for up to 90 million stars with better than 5% accuracy, and measure proper motions with an accuracy approaching microarcsec per year. If hierarchical CDM is correct, there should be thousands of coherent streamers that make up the outer halo, and hundreds of partially phase-mixed structures within the inner halo. A satellite experiencing the disk potential no longer conserves its angular momentum and its orbit plane undergoes strong precession (see Figure 10c,d). In Figure 10a,b, Helmi et al. (1999) demonstrate the relative ease with which Gaia will identify substructure within the stellar halo.