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The Chemical Impact of Stellar Mass Loss

The Chemical Impact of Stellar Mass Loss. Rosemary Wyse Johns Hopkins University. Gerry Gilmore, John Norris, Mark Wilkinson, Vasily Belokurov, Sergei Koposov, Wyn Evans, Dan Zucker, Anna Frebel, David Yong. Elemental abundances. Field stars and dwarf spheroidal members

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The Chemical Impact of Stellar Mass Loss

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  1. The Chemical Impact of Stellar Mass Loss Rosemary Wyse Johns Hopkins University Gerry Gilmore, John Norris, Mark Wilkinson, Vasily Belokurov, Sergei Koposov, Wyn Evans, Dan Zucker, Anna Frebel, David Yong

  2. Elemental abundances • Field stars and dwarf spheroidal members • Massive-star mass function (core-collapse SNe) • Invariant • Mixing in interstellar medium • Surprisingly efficient • Carbon-rich (single) stars at very low [Fe/H] • But also carbon-normal ultra-metal-poor stars

  3. Elemental Abundances: beyond metallicity < • Core collapse supernovae have progenitors > 8 M and explode on timescales ~ 107 yr, less than typical duration of star formation • Main site of -elements, e.g. O, Mg, Ti, Ca, Si • Low mass stars enriched by only Type II SNe show enhanced ratio of -elements to iron, with value dependent on mass distribution of SNe progenitors – if well-mixed system, see IMF-average • Type Ia SNe produce very significant iron, on longer timescales, few x 108 – several 1010yr (WD in binaries) after birth of progenitor stars

  4. Schematic [O/Fe] vs [Fe/H] Wyse & Gilmore 1993 IMF biased to most massive stars Type II only Plus Type Ia Fast Slow enrichment Low SFR, winds.. Self-enriched star forming region. Assume good mixing so IMF-average yields

  5. IMF dependence due to different nucleosynthetic yields of Type II progenitors of different masses Salpeter IMF (all progenitor masses) gives [/Fe] ~ 0.4; Change of IMF slope of ~1 gives change in [/Fe] ~ +0.3 (Wyse & Gilmore 92) Ejecta Kobayashi et al 2006 Progenitor mass Gibson 1998

  6. Elemental abundances in old stars Ruchti et al 2011,12 Thick disk extends to -2 dex, same enhanced [α/Fe] as halo stars same massive-star IMF, short duration of star formation little scatter – fixed IMF, good mixing, down to [Fe/H] < -3 dex Stars from RAVE survey, candidate metal-poor disk, follow-up echelle data

  7. Bulge Matches Thick Disksame massive-star IMF Gonzales et al 2011

  8. Extended, low-rate star formation and slow enrichment with gas retention, leads to expectation of ~solar (or below) ratios of [/Fe] at low [Fe/H], such as in LMC stars Hiatus then burst Pompeia et al 2008 Smith et al 2003 LMC: solid Local disk Gilmore & Wyse 1991

  9. Gilmore et al; Norris et al 10 BooI Simon et al 10 Leo IV Boo I  dSphs vs. MWG abundances: SFH(from A. Koch, 2009 + updates) Frebel et al 10 Scl Leo IV ◊ Scl        Shetrone et al. (2001, 2003): 5 dSphs Koch et al. (2008): Hercules Sadakane et al. (2004): Ursa Minor Shetrone et al. (2008): Leo II Monaco et al. (2005): Sagittarius Frebel et al. (2009): Coma Ber, Ursa Major Koch et al. (2006, 2007): Carina Aoki et al. (2009): Sextans Letarte (2006): Fornax Hill et al. (in prep): Sculptor

  10. Same ‘plateau’ in [α/Fe] in all systems at lowest metallicities • Type II enrichment only: massive-star IMF invariant, and apparently well-sampled/mixed • Stellar halo could form from any system(s) in which star-formation is short-lived, and inefficient so that mean metallicity kept low • Star clusters, galaxies, transient structures… • Complementary, independent age information that bulk of halo stars are OLD further constrains progenitors (e.g. Unavane, Wyse & Gilmore 1996)

  11. Star Counts: Invariant Low-Mass IMF Main sequence luminosity functions of UMi dSph and of globular clusters are indistinguishable. HST star counts Wyse et al 2002 UMi dSph stars have narrow range of ages, all old M92  M15  0.3M

  12. Low-Mass Stellar MF in Bulge: Zoccali et al 2000 (M15) Matches local disk (Kroupa 2000) And M15 – which matches the UMi dSph: Low-mass IMF invariant wrt metallicity, time..

  13. Carina dSph CMD Stetson et al 2011 Very extended, non-monotonic star formation history

  14. Carina dSph – extended, bursty star formation history Carina data: bursts + inhomogeneous star formation Massive star IMF invariant Koch et al inc RW 2008

  15. Age estimates: younger indeed higher [α/Fe] Lemasle et al 2012

  16. A much simpler system: Bootes I ‘ultra-faint’ dwarf M* ~ 4 x 104 M, dist ~ 65 kpc SDSS Discovery CMD (Belokurov et al, inc RW, 2006b) Subaru (Okamoto, PhD, 2010)

  17. [Fe/H] distributions and radial dependence Norris, RW et al 2010 16 stars Dwarf spheroidal galaxies have well-defined peaks, with low metallicity tails: self-enriched, from primordial gas? Then lost most baryons – M/L high. Segue 1, 7 stars Very luminous globular cluster  lacks low-metallicity tail; most clusters do not self-enrich in Fe; Need enough compact baryons

  18. Alpha Abundances: • 8 stars in Boo I, VLT UVES • Double-blind analysis (Gilmore et al 2012)  minimal scatter Boo-119 is CEMP-no star; open dots are field CEMP CEMP-no star Segue 1-7 has [Mg/Fe] ~ 0.94 (Norris et al 2010)

  19. Carbon-enhanced star in Segue 1 (triangles) and BooI (circles) Norris et al 2010a,b No s-process plus high [Mg/Fe]

  20. Including data for Boo I stars from Lai et al 2011

  21.            Two modes of enrichment? • Unmixed, very early, enriched by individual supernovae from zero-metal stars? • Extremely well mixed, fully sample massive-star IMF – minimal scatter in element ratios • Boo I probably lost 90% of baryons – metals? [Fe/H] time ISM mixing scale

  22. Conclusions • Lack of variations in elemental abundances probably produced by core-collapse supernovae argue for invariant massive-star IMF • Star counts imply fixed low-mass IMF • Overall patterns determined by star-formation history • Small scatter implies well-mixed ISM • WHY? And HOW?

  23. Large Scale Flows • Chemical evolution plus global star formation rates argue for gas replenishment • High velocity clouds exist • Galactic Fountain • Cold Flows from Cosmic Web • Accretion from satellite galaxies (Magellanic Stream) • Radial migration

  24. Boötes I ~ • MV ~ -6.3, M* ~ 4 x 104 M (Kroupa IMF), distance of ~ 65kpc, half-light radius ~ 250pc (< dark matter scalelength?), central velocity dispersion ~ 3-6 km/s (?), derived (extrapolated) mass within half-light radius ~ 106-7 M, M/L ~ 103, mean dark matter density ~ 0.1M/pc3 collapse at redshift z > 10 • Color-magnitude diagram consistent with old, metal-poor population, similar to classic halo globular cluster • More luminous dSph have very varied SFHs ~ Belokurov et al 06; Gilmore et al 07; Martin et al 08; Walker et al 09; Okamoto et al 10

  25. Koposov, et al (inc RW), 2011b • Getting the most from Flames on VLT: Bootes I field, • ~1 half light radius FOV, 130 fibres, 12 x 45min integrations • Repeated observations allow detection of variability: • 110 non-variable (giant) stars (< 1km/s) Analyse spectra in pixel space; Retain full covariance: map model spectra onto data, find ‘best’ match values of stellar parameters (gravity, metallicity, surface temperature) with a Bayesian classifier. Black: data r=19; red=model

  26. Koposov, et al (inc RW), 2011b Identify 37 members, based on line-of-sight velocity, metallicity and stellar gravity (should be giants, dwarfs will be foreground field halo stars) Previous literature value (already reduced)

  27. Field of Streams(and dots) Belokurov et al (inc RW, 2006)  Segue 1  Boo I Outer stellar halo is lumpy: but only ~15% by mass (total mass ~ 109M) and dominated by Sgr dSph stream SDSS data, 19< r< 22, g-r < 0.4 colour-coded by mag (distance), blue (~10kpc), green, red (~30kpc)

  28. Norris, RW et al 2010 Wide-area spectroscopy Red: Segue 1 Black: Boo I Geha et al І • Members well beyond the nominal half-light radius in both • Stars more iron-poor than -3 dex (10-3 solar) exist in both • Extremely rare in field halo, membership very likely • Very far out, parameters and velocity confirmed by follow-up: Segue 1 is very extended! (as is Boo I) • Both systems show a large spread in iron • Implies dark halo for self-enrichment (cf Simon et al 2011, 6 stars in Segue 1, 7 in total)

  29. Wyse & Gilmore 1992 Salpeter IMF slope: -1.35 Scalo: -1.5 Matteucci for Bulge: -1.1

  30. Chemical Abundances: Boo I & Segue 1 • Spectroscopic surveys with the 2dF/AAΩ fibre-fed MOS; stars selected from SDSS to follow discovery CMD: wide-area mapping unique capability of 2dF • 400 fibres, 2-degree FOV, dual beam, chemical abundances from blue spectra,R ~5000, range 3850-4540Å. Membership based on radial velocity (to better than 10 km/s) and the derived values of stellar parameters • Iron from calibration of Ca II K line (3933Å, as field halo surveys, Beers et al 99), +/- 0.2 dex (Norris et al 08) • Carbon from synthesis of CH G-band, +/- 0.2 in Boo I and +/- 0.4 in Segue 1 (Norris et al 10) • Follow-up UVES echelle data, [Fe/H] +/- 0.1dex; [C/Fe] for 1 star

  31. Caveat: Segue 1 in very complex part of the Galaxy • Very flat (bimodal?) metallicity distribution, unlike other dwarf galaxies: contamination? • Extended structure around it • Same distance and line-of-sight as Sgr stream, but different velocity (Niederste-Ostholt et al wrong orbit for Sgr stream) • Distance and velocity, line-of-sight match Orphan stream (Newberg et al 2010, Koposov et al inc RW 2011a) • What is the `300km/s stream’? • Extremely wide-field mapping needed to be assured of status

  32. Segue 1 • MV ~ -1.5, M* ~ 600 M, distance of ~25kpc, half-light radius ~ 30pc (?), velocity dispersion ~ 4 km/s (?), derived mass within half-light radius ~ 3 x 105M(?), M/L ~ 2000 (?), again <ρ>DM ~ 0.1 M/pc3 and high collapse redshift • Superposed on Sgr tidal debris, close in distance and velocity (?), contamination likely (Niederste-Ostholt et al 09); unlikely (Simon et al 2010) • Again old, metal-poor population Belokurov et al. 07; Martin et al 08; Geha et al 09; Walker et al 09; Simon et al 2010

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