KIAA Lectures Beijing, July 2010 Ken Freeman, RSAA Lecture 3: the Galactic bulge and the globular clusters

NGC 4594 : a classical r 1/4 bulge

NGC 4565 : a boxy bulge

Our Galaxy has a small boxy bar-bulge

NGC 5907: no bulge at all

From lecture 1 …. Forming large galaxies with small or no bulges is currently difficult in CDM because of the relatively active ongoing merger history. Establishing the merger history of the Milky Way observationally is a major goal for Galactic Archaeology. We need to understanding how the Galactic bulge formed : is it even partly a merger product or did it form entirely through internal processes (eg disk and bar instability)

Large classical bulges as in the Sombrero galaxy are believed to be merger products. Merger dynamics and violent relaxation leads to bulges with the characteristic r 1/4 -law light distribution (Sersic index ~ 4). Classical bulges are common in early type galaxies but become progressively rarer towards later types. They share some structural, dynamical and population properties with the lower- luminosity ellipticals Different kinds of bulges

Later type galaxies like the Milky Way mostly have small near-exponential boxy bulges, rather than r 1/4 bulges. (eg Courteau et al 1996) observations: Kuijken & Merrifield 1995, Bureau & KF 1999 Chung et al theory: Combes & Sanders Boxy bulges, as in our Galaxy, are associated with bars, believed to form via bar-buckling instability of disk. They are probably not merger products

NGC 5746: gas kinematics in a boxy bulge show the signature of orbits in a bar potential (Bureau & Freeman 1999) NGC 5746 [NII] 6584Å HH [NII] 6548Å

Bar-buckling to form a boxy/peanut bulge The disk suffers a bar instability : very common for fairly cold disks The bar buckles vertically, driven by horizontal and vertical resonances, and forms a boxy/peanut bulge: it takes a few bar revolutions to make this instability go (Combes et al 1990, Athanassoula et al 2007). The whole process takes 2-3 Gyr after the formation of the disk The rotation of boxy bulges is cylindrical: i.e. V rot only weakly dependent on height above the plane structure velocity field

The maximum vertical extent of peanuts occurs near radius where vertical and horizontal Lindblad resonances occur ie where  b =  -  /2 =  - z /2 (remember: both  and z depend on the amplitude of the oscillation) Stars in this zone oscillate on orbits which support the peanut shape. Orbits supporting the peanut

So far, seen classical bulges, probably merger products boxy/peanut bulges, probably disk instability products There is a third kind of bulge-like structure which looks like an enhancement of the surface brightness profile above the exponential disk but appears to lie in the disk, from its shape and kinematics. These are the pseudobulges (Kormendy 1993): they are believed to be generated by secular processes associated with the angular momentum transport by bars or weakly oval disks. They often show active star formation within the pseudobulge region. (Recall that bars are very common : about 2/3 of disk galaxies show some kind of bar structure in NIR images)

M83 in blue light (L) and K light (near-IR) (R) The bar is much more obvious in the near-IR. The bar extends well beyond the central bulge.

Kormendy & Kennicutt 2004 NGC 6384 pseudobulge

Another example of a starforming pseudobulge (HST) M. Carollo et al 1998

The kinematics of pseudobulges : V/  above oblate curve

The terms pseudobulge and secular evolution have become a bit mis-used. I think pseudobulge is best reserved for these flat enhancements that look like bulges only in their surface brightness profiles. There is nothing pseudo about the Galactic bulge. Secular evolution means slow relative to the dynamical time, like the slow transport of matter into the central regions via torques from a bar or oval disk. There is nothing secular about the bar-buckling scenario.

Kinematics of classical vs boxy bulges Falcon-Barosso et al 2004 NGC 5866 NGC 7332

Kinematics of classical bulge (NGC 5866): non-cylindrical rotation (SAURON) Falcon-Barosso et al 2004

Kinematics of boxy bulge (NGC7332): near-cylindrical rotation (SAURON) Falcon-Barosso et al 2004

The Galactic Bar- Bulge small exponential bulge - typical of later-type galaxies. Launhardt 2002

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

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

Near the center of the bar/bulge is a younger population, on scale of about 100 pc : the nuclear stellar disk (M ~ 1.5 x 10 9 M_sun) and nuclear stellar cluster (~ 2 x 10 7 M_sun ) in central ~ 30 pc. (Launhardt et al 2002) ~ 70% of the luminosity comes from young main sequence stars.

Later type galaxies like the Milky Way mostly have small near-exponential boxy bulges, rather than r 1/4 bulges. (eg Courteau et al 1996) These small boxy bulges are probably not merger products: more likely generated by bar-buckling instability of disk. We might then expect some similarities of stellar population between the bulge and the surrounding disk and thick disk How did the Galactic Bulge form ?

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

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

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

bulge thick disk thin disk Meléndez et al 2008 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.

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 bar-forming and bar-buckling process takes 2-3 Gyr to act after the disk settles 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 galactic bulge is rotating, like most other bulges: Rotation (Beaulieu et al 2000) K giants from several sources and planetary nebulae (+) Velocity dispersion of inner disk and bulge are fairly similar not easy to separate inner disk and bulge kinematically Bulge ends at |l| ~ 12 o

As expected for exponential disk in R and z : scaleheight ~ 300 pc, scalelength 3-4 kpc. Velocity dispersion increases from ~ 15 km/s at 18 kpc to ~ 100 km/s near the center (similar to bulge). This makes it difficult to separate disk and bulge stars kinematically Lewis & KCF R (kpc) log (velocity dispersion) Velocity dispersion of the thin disk

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

N-body model seen from galactic pole

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

The kinematics of the model are as observed for boxy bulges: cylindrical rotation b = 0.5 b = 9.5 Detailed velocity data not yet available for the galactic bar/bulge: survey in progress. Model fits well to limited data available now

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

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

The ARGOS bulge survey We are doing a spectroscopic survey of the bulge with the AAT and AA  to determine whether its kinematics are consistent with the bar-buckling scenario and to derive limits on any underlying classical bulge. (Melissa Ness, KCF et al) Observe at Ca triplet ~ 8600 Å, resolution = 13,000, SN ~ 70 Magnitudes chosen to cover entire sightline through the bulge 28 fields of 1000 stars each, in bulge and surrounding thin and thick disk (have spectra of 23,000 stars so far) sun

Sample size sufficient to detect 5% merger generated bulge underlying an instability bulge. Selected stars in each field from 2MASS Mainly red clump giants along line of sight M K = -1.6, (J-K) 0 = 0.65 Colour cuts determined using Schegel reddening in each field Selection criteria do not exclude metal-poor stars: expect to find stars of the inner stellar halo. Also … In CDM cosmology formation scenarios, the first stars will be concentrated in the bulge region Sample Selection Criteria: 28,000 stars in 28 Fields

Spectrum observed at Ca triplet with AAOmega See lines of Fe, Al, Ca, Ti, Si, Mg, O

Rotation Curves for 4 Fields of Latitude: From output velocities of ~ 23, 000 stars (error < 1.2km/s) V gc = V hc + 220sin(l )cos(b) [sin(b)sin(25) + cos(b)cos(25)cos(l − 53)] cylindrical rotation [Athanassoula] near sidebulge far side

Find the expected metal-poor halo stars in the bulge region. They do not rotate as fast as the more metal-rich stars of the bulge (previously described by Paul Harding 1993) Are they just the stars of the inner halo, or are they the first stars, or is there no difference ?

The metal-poor stars in the bulge region rotate more slowly than the metal-rich stars: they probably belong to the inner Galactic halo

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. First stars are in orbits of fairly high eccentricity, rather similar to observed eccentricity distribution for metal-poor stars in the galactic halo

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

Distribution in present galaxy of debris from peaks selected at z > 12 (Moore et al 2006). Dashed cuve shows slope for metal-poor halo.

The bulge is not a dominant feature of our Galaxy - only about 20% of the light. The bulge is probably an evolutionary structure of the disk, rather than a feature of galaxy formation in the early universe. Structure and kinematics (so far) are well represented by product of disk instability. The  -enhancement indicates that star formation in this inner disk/bulge region proceeded rapidly. The bulge structure may be a few Gyr younger than its stars. The Galactic Bulge - summary

The M31 bulge peak V rot ~ 100 km/s  (0) ~ 140 km/s Simien et al 1979, McElroy 1983 Rotation and velocity dispersion of its bulge are slightly larger than for the MW bulge

Athanassoula et al (2006) Beaton et al (2007) : J-band isophotes M31 has a classical bulge plus an inner boxy bar/bulge, from detailed comparison of the isophote structure with their N-body models of boxy bar/bulges. The flat part of the bar extends about 1.4 times further in radius than the boxy bulge J 2MASS J image

Globular clusters The globular clusters in our Galaxy are all old. Their [Fe/H] distribution shows two modes: bulge/disk clusters with [Fe/H] > -1 halo clusters with [Fe/H] < -1

Recent version of Zinn's two-population figure, from Cote (1999) Metal rich clusters R < 4 kpc V rot = 157 km/s  = 70 km/s thick disk kinematics Metal poor clusters R > 4 kpc V rot = 22 km/s  = 123 km/s halo kinematics

Zinn (1985) established: The clusters more metal-rich than [Fe/H] ≈ -1 form a disk system with a highly flattened spatial distribution and a significant rotation velocity. The scale height and rotational velocity of the system is comparable to that of the thick disk. (†) The clusters more metal-poor than [Fe/H] ≈ -1 are part of the halo population - an essentially spherical distribution about the Galactic center, with a small rotation velocity and a large velocity dispersion. (†) Minniti (1995) argued that the metal-rich globular clusters near the Galactic Centre were more likely associated with the bulge rather than the thick disk.

Globular clusters are potentially very important in galactic archaeology, but we do not understand how they form and therefore what they represent in the process of galaxy formation. Some appear to be associated with Sgr, Mon, Can Maj streams so probably formed in the parent objects which were accreted The old globular clusters in the Galaxy, the LMC and Fornax are all coeval within 1 Gyr. Most clusters except  Cen are chemically homogeneous in heavy elements (Ca, Fe … ), though not in light elements like Na and O

ACS survey of Galactic Globular Clusters: new relative ages Coeval clusters over whole [Fe/H] range, plus younger clusters with age-metallicity relation - ~12 Gyr

 The Old Halo clusters show a weak gradient: R gc ≤ 6 kpc  [Fe/H]  = (42 clusters) 6 ≤ R gc ≤ 15 kpc  [Fe/H]  = (17 clusters) R gc ≥ 15 kpc  [Fe/H]  = (11 clusters) Metallicity Gradients in the Galaxy  The Bulge/Disk clusters show no abundance gradient (but these clusters occur only in the central regions).  The Young Halo clusters show no gradient : R gc ≤ 15 kpc  [Fe/H]  = (14 clusters) R gc ≥ 15 kpc  [Fe/H]  = (16 clusters)

In the LMC and M33, some of the globular clusters are very old and metal-poor like the clusters in the Galaxy - about 12 Gyr old. but some are very young - only a few million years old. Why is this important ? We would like to know how these very dense clusters form. The LMC and M33 are able to form them now, but our Galaxy is not : what is the difference ? and others have intermediate ages - 10 million years to a few Gyr

The LMC is forming globular clusters now: it has globular clusters of all ages. NGC 1850 M = 6 x 10 4 M  Age = 90 Myr

As in the LMC, there are young and old clusters, but 85% of the old M33 clusters lie in a kinematic halo, whereas the old clusters in the LMC are all moving in the disk Chandar et al 2002 Globular clusters in M33

Globular Clusters: multiple main sequences and abundance anomalies eg NGC 2808 has three main sequences, believed to represent three discrete He levels up to Y ~ 0.40 Pre-enrichment must generate discrete levels of He and must not affect [Fe/H] and [  /Fe]. Requires pollution by high- temperature H-burning in previous generation of stars. Important for understanding how the clusters form.

The Na-O anticorrelation Most clusters are homogeneous in the heavy elements (Ca …) but not in light elements. The Na/O anticorrelation is seen only in globular clusters, not in the field stars, so it is somehow related to the globular cluster environment or pre-environment. It requires enrichment from a previous generation of massive stars but must not affect [Fe/H] and [  /Fe] It should be easy to find the debris of disrupted globular clusters by chemical tagging: they will show the Na/O anticorrelation

The surrounding galaxy environment can provide multiple generations of enrichment. Need to get material into nucleus : likely to be a sporadic dynamically driven process, delivering discrete levels of enrichment at a few particular times. Nuclei of low-luminosity bulgeless spirals are much like massive GCs in velocity dispersion, mass, surface density, subsolar metallicity. UVES spectra of spiral nuclei indicate continuing episodic star formation (Walcher et al 2006) Globular clusters and galactic nuclei Large and inhomogeneous clusters like  Cen may be surviving nuclei of accreted and stripped galaxies.

Histograms of J z for Gratton's (2003) sample of nearby metal- poor stars with well-measured chemical abundances The retrograde  Cen feature is probably associated with the accretion event that brought  Cen into the MW. Its stars have chemical anomalies like those of  Cen itself (Wylie et al 2010) Meza et al 2006 Debris from the  Cen accretion event

Conclusion on globular clusters Globular clusters are potentially useful tracers of the early phases of galaxy evolution but we don't yet Understand what physical conditions they trace. What conditions are needed for their formation ? We see globular cluster formation in several different situations now: merging galaxies, starburst galaxies, the disks of M33 and LMC. Yet in the large spirals like the MW and M31, the globular clusters are mostly very old: the oldest clusters have very similar ages throughout the Local Group No clear picture emerges about the properties of globular clusters in different galaxies. Chemical evolution from the pre-cluster environment may be important for understanding chemical anomalies We don't really understand what globular clusters represent in the context of galaxy formation.