More Clues to Galaxy Formation: Massive Globular Clusters, Stochastic Self-Enrichment, and Mass/Metallicity Correlations. NGC 4696. HST/ACS. Harris && 2006. Pregalactic dwarf. Proto-GCs. Young massive star clusters (YMCs) forming at ~10 5 M 0 in starburst dwarfs today.
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More Clues to Galaxy Formation: Massive Globular Clusters, Stochastic Self-Enrichment, and Mass/Metallicity Correlations
Harris && 2006
Young massive star clusters (YMCs) forming at ~105 M0 in starburst dwarfs today
Starburst dwarf NGC 5253 (ESO/HST)
Harris && 2008
Harris && 2006
Bimodal or not?
Harris && 2008
Harris && 2006
Is this effect caused by ---
(1) A gradual shift of the blue sequence to redder color at higher luminosity? (Mass/Metallicity relation)
(2) The disappearance of bimodality altogether at the highest masses? (Threshold enrichment effect)
(3) An artifact of photometric measurement procedures? (i.e. not real)
If it’s a true, physical MMR then Z ~ M1/2 at high mass, and it may smoothly connect upward to the UCD regime.
Does it continue to low mass? Why no red-sequence MMR?
Is it present in all galaxies? What is its astrophysical origin?
- Multiple populations within a single GC
- Different scaling of size vs. mass
Evstigneeva et al. 2008
The systematic properties of globular clusters begin to change for M > 2 x 106 M0 …
- Appearance of the MMR
Can be helped (partially) by constructing composite samples; e.g grouping Virgo Cluster Survey galaxies into 4 luminosity groups (Mieske && 2006) or combining several supergiants (Harris && 2006)
But if amplitude of MMR differs from one galaxy to another, net effect will be diluted in composite samples
The basic feature of bimodality is a first-order and (probably) universal effect. The MMR is a second-order effect and harder to trace. Though new, much confusion already exists:
Category 1:MMR is present and measurable
M87, NGC 1399, several other BCGs and gE’s
Category 2:MMR is not present
M49; any others?
Category 3:presence of MMR not decidable; GC sample too small or does not extend to high enough luminosity
Milky Way; M31; dwarf galaxies; most spirals; GC-poor E’s
1: strong MMR
2: no MMR
3: Not decidable
Milky Way GCs
NGC 5128: d=3.8 Mpc
Globular clusters are easily resolved at <1’’ seeing
Photometry must account for individually different scale sizes
GC profile as seen on image =
PSF Intrinsic GC profile
rh ~ 1 – 5 parsecs; averages 3 pc 0.3” width
d = 50 Mpc
2 rh ~ 6 pc 0.025”
fwhm(PSF) = 0.5”
starlike! psf-fitting photometry is fine
Several regimes determined by distance; no single photometric method is suitable for all regimes
Gemini-S + GMOS, Wehner & Harris
HST/ACS imaging of GCs around 6 central supergiants in Abell-type clusters (Harris et al. 2006, 2008)
(B,I) bandpasses metallicity-sensitive
Thousands of GCs per galaxy, thus good statistical samples and big luminosity (mass) range
4 distinguishable regimes: compare fwhm of stellar PSF with intrinsic cluster size D (= 2 rh),half-light diameter
Well resolved: D >> fwhm(PSF)
Partially resolved: D ~ fwhm
Marginally resolved: D ~ 0.1 – 0.3 fwhm
Unresolved (starlike): D < 0.1 fwhm
All this is subject to S/N considerations …
NGC 1407 Eridanus d=23 Mpc MV = -22.35
NGC 3258 Antlia 41 Mpc -21.87
NGC 3268 Antlia 41 Mpc -21.96
NGC 3348 CfA69 41 Mpc -22.13
NGC 4696 Centaurus 42 Mpc -23.31
NGC 7626 Pegasus I 49 Mpc -22.58
(M87 Virgo 16 Mpc -22.4)
(Partial list – biggest GCSs out of 12 studied)
D = 6 pc at d ~ 40 Mpc half-light profile width ~ 0.03”
compare PSF fwhm = 0.1” marginally resolved
ISHAPE sample fits
1 px = 0.05”
fwhm a=1.3 px
b/a = 0.91
fwhm a=0.82 px
b/a = 0.50
Simulations show that the systematically correct intrinsic D (FWHM of GC profile) is returned for D > 0.1 (PSF) (transition boundary from unresolved to marginally resolved)
Growth curves for simulated GC profiles convolved with PSF
ISHAPE solve for best-fit D
Measure magnitude through 2.5-px aperture, corrected back to the growth curve for a starlike profile
More tests …
Measured size a not affected by modestly elliptical shape
b/a, q returned correctly for a > 0.1 PSF
Previous PSF-fitting data (Harris && 2006)
N=12000 brighter than MI = -8. Largest sample in existence!
Blue sequence: gradual changeover to MMR toward higher mass
Z ~ M0.3+-0.1
RMIX fits of bimodal gaussians within selected magnitude intervals: forces two modes into the solution, but (a) less affected by field contamination, (b) avoids the strong assumption imposed by a ‘linear fit’
Trends (?) versus galactocentric distance and metallicity: projection effects, or intrinsic?
Low-metallicity GCs average larger size at any galactocentric zone
SNe from >8 M0 stars enrich lower-mass stars while still in formation
Salpeter IMF 0.3 100 M0 and SF efficiency f* ~ 0.3
Woosley/Weaver SN yields, and fraction fZ of heavy elements retained in GMC
What is responsible for the metallicity distribution function (MDF)?
Bailin & Harris 2008
Is a proto-GC
- PRE-enriched from the surrounding GMC gas?
- internally SELF-enriched by its own SNe within the first few Myr?
- stochastic? (can self-enrichment be responsible for the internal dispersion of the MDF?)
Internal dispersion of MDF due to statistical variation in NSN
NSN ~ 1 per 100 M0
Stochastic self-enrichment fails to explain the MDF dispersion at any cluster mass higher than 104 M0
Proto-GC = truncated isothermal sphere logarithmic potential F(R). All SNe go off while PGC is still highly gaseous; all ejected energy absorbed and thermalized.
Gas will leave if outside an “escape radius” defined by total energy > potential energy at edge of cloud.
Ejecta become efficiently retained at a characteristic mass (after star formation)
Two additional, major factors to add:
- reff ~ M1/2 at high mass
- fZ is a strong function of M(init) and thus reff as well
Combined effects of pre-enrichment, self-enrichment, and mass/radius relation