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More Clues to Galaxy Formation: Massive Globular Clusters, Stochastic Self-Enrichment, and Mass/Metallicity Correlations NGC 4696 HST/ACS Harris && 2006.

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Presentation on theme: "More Clues to Galaxy Formation: Massive Globular Clusters, Stochastic Self-Enrichment, and Mass/Metallicity Correlations NGC 4696 HST/ACS Harris && 2006."— Presentation transcript:

1 More Clues to Galaxy Formation: Massive Globular Clusters, Stochastic Self-Enrichment, and Mass/Metallicity Correlations NGC 4696 HST/ACS Harris && 2006

2 Starburst dwarf NGC 5253 (ESO/HST) Pregalactic dwarf Proto-GCs Young massive star clusters (YMCs) forming at ~10 5 M 0 in starburst dwarfs today

3 Harris && 2008Harris && 2006 Bimodal or not?

4 Harris && 2008Harris && 2006 Bimodal or not? Disagreements ahead --

5 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) Serious questions persist! If it’s a true, physical MMR then Z ~ M 1/2 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?

6 The systematic properties of globular clusters begin to change for M > 2 x 10 6 M 0 … - Appearance of the MMR - Multiple populations within a single GC  Cen (Villanova && 2007) - Different scaling of size vs. mass Evstigneeva et al. 2008

7 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 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: 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

8 1: strong MMR2: no MMR3: Not decidable [Fe/H] MVMV Milky Way GCs Most galaxies do not have clusters in the 10 6 – 10 7 M 0 range

9 First, let’s get the measurements straightened out. 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 r h ~ 1 – 5 parsecs; averages 3 pc  0.3” width

10 NGC 3311/3309 (A1060) d = 50 Mpc 2 r h ~ 6 pc  0.025” fwhm(PSF) = 0.5”  starlike! psf-fitting photometry is fine Gemini-S + GMOS, Wehner & Harris Several regimes determined by distance; no single photometric method is suitable for all regimes

11 4 distinguishable regimes: compare fwhm of stellar PSF with intrinsic cluster size D (= 2 r h ), 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 Aperture photometry r(ap) adjusted for D PSF-fitting photometry All this is subject to S/N considerations … 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

12 NGC 1407 Eridanus d=23 Mpc M V = NGC 3258 Antlia 41 Mpc NGC 3268 Antlia 41 Mpc NGC 3348 CfA69 41 Mpc NGC 4696 Centaurus 42 Mpc NGC 7626 Pegasus I 49 Mpc (M87 Virgo 16 Mpc -22.4) D = 6 pc at d ~ 40 Mpc  half-light profile width ~ 0.03” compare PSF fwhm = 0.1”  marginally resolved HST/ACS Imaging program for BCGs (Partial list – biggest GCSs out of 12 studied)

13 Photometric technique: - Uniform catalog of detected objects with DAOPHOT - Construct PSF from average of many bright starlike objects - For each individual source, convolve PSF with “King30” model GC profile and vary D(model) to obtain best match (ISHAPE; Larsen 1999) - finally, use fixed-aperture photometry corrected for profile width to obtain final magnitude in each band S/N=441 fwhm a=1.3 px b/a = 0.91 S/N=24 fwhm a=0.82 px b/a = 0.50 S/N=108 starlike ISHAPE sample fits 1 px = 0.05” HST/ACS

14 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 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) Our regime

15 More tests … Measured size a not affected by modestly elliptical shape b/a,  returned correctly for a > 0.1 PSF S/N > 50 !!

16 Full, profile-corrected aperture photometry for 6 supergiant ellipticals Previous PSF-fitting data (Harris && 2006) Trend lines: - blue and red? - linear slope? or top end only? - how steep? N=12000 brighter than M I = -8. Largest sample in existence!

17 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’ Red sequence: no trend Blue sequence: gradual changeover to MMR toward higher mass Z ~ M

18 The top end: uni- or bi-modal?

19 A detour: the measured cluster sizes

20 Trends (?) versus galactocentric distance and metallicity: projection effects, or intrinsic? Low-metallicity GCs average larger size at any galactocentric zone

21 The MMR is not due to an unaccounted-for size-mass relation.

22 What is responsible for the metallicity distribution function (MDF)? 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?) Input assumptions to self-enrichment model: SNe from >8 M 0 stars enrich lower-mass stars while still in formation Salpeter IMF 0.3  100 M 0 and SF efficiency f * ~ 0.3 Woosley/Weaver SN yields, and fraction f Z of heavy elements retained in GMC and thus Bailin & Harris 2008

23 N SN ~ 1 per 100 M 0 Pre-enrichment level for f Z = 0.08 Internal dispersion of MDF due to statistical variation in N SN Stochastic self-enrichment fails to explain the MDF dispersion at any cluster mass higher than 10 4 M 0

24 Two additional, major factors to add: - r eff ~ M 1/2 at high mass - f Z is a strong function of M(init) and thus r eff as well Proto-GC = truncated isothermal sphere  logarithmic potential  (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)

25 Combined effects of pre- enrichment, self-enrichment, and mass/radius relation Match to BCG data for 6 galaxies -pre-enrichment of each “mode” (blue, red) tuned to match mean color -self-enrichment drives shape of mean MDF at high mass

26 Basic features of the model: - No MMR for cluster masses < ~10 6 M 0 (i.e., sequences vertical) - Very metal-poor, very massive GCs should be rare (anywhere) - blue and red sequence converge at high-mass end - Similar red-sequence MMR should exist at top end, but smaller amplitude - Internal dispersion and mean metallicity of each mode driven by pre-enrichment Wehner && 2008

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