The Chemical Impact of The Chemical Impact of Stellar Mass LossStellar Mass Loss

Rosemary WyseRosemary Wyse

Johns Hopkins UniversityJohns Hopkins University

Gerry Gilmore, John Norris, Mark Wilkinson, Vasily Belokurov, Sergei Koposov, Wyn Evans, Dan Zucker, Anna Frebel, David Yong

Elemental abundancesElemental abundances

Field stars and dwarf spheroidal membersField stars and dwarf spheroidal members Massive-star mass function (core-collapse Massive-star mass function (core-collapse

SNe)SNe) Invariant Invariant

Mixing in interstellar mediumMixing in interstellar medium Surprisingly efficientSurprisingly efficient

Carbon-rich (single) stars at very low Carbon-rich (single) stars at very low [Fe/H][Fe/H] But also carbon-normal ultra-metal-poor starsBut also carbon-normal ultra-metal-poor stars

Elemental Abundances: Elemental Abundances: beyond metallicitybeyond metallicity

Core collapse supernovae have progenitors > 8 MCore collapse supernovae have progenitors > 8 M

and explode on timescales ~ 10and explode on timescales ~ 107 7 yr, less than yr, less than typical duration of star formationtypical duration of star formation

Main site of Main site of -elements, e.g. O, Mg, Ti, Ca, Si-elements, e.g. O, Mg, Ti, Ca, Si Low mass stars enriched by only Type II SNe show Low mass stars enriched by only Type II SNe show

enhanced ratio of enhanced ratio of -elements to iron, with value -elements to iron, with value dependent on mass distribution of SNe progenitors dependent on mass distribution of SNe progenitors – if well-mixed system, see IMF-average– if well-mixed system, see IMF-average

Type Ia SNe produce very significant iron, on Type Ia SNe produce very significant iron, on longer timescales, few x 10longer timescales, few x 1088 – several 10 – several 101010yr (WD in yr (WD in binaries) after birth of progenitor starsbinaries) after birth of progenitor stars

<

Schematic [O/Fe] vs [Fe/H]Schematic [O/Fe] vs [Fe/H]Wyse & Gilmore 1993

Slow enrichmentLow SFR, winds..

Fast

IMF biased to most massive stars

Self-enriched star forming region.Assume good mixing so IMF-average yields

Type II onlyPlus Type Ia

Gibson 1998

Progenitor mass

Eje

cta

Salpeter IMF(all progenitor masses) gives [/Fe] ~ 0.4;Change of IMF slope of ~1 giveschange in [ /Fe]~ +0.3 (Wyse & Gilmore 92)

IMF dependence due to different nucleosynthetic yields of Type II progenitors of different masses

Kobayashi et al 2006

Elemental abundances in old Elemental abundances in old starsstars

Ruchti et al 2011,12

Thick disk extends to -2 dex, same enhanced [α/Fe] as halo starssame massive-star IMF, same massive-star IMF, short duration of star short duration of star formationformationlittle scatter – fixed IMF, little scatter – fixed IMF, good mixing, down to [Fe/H] good mixing, down to [Fe/H] < -3 dex< -3 dex

Stars from RAVE survey, Stars from RAVE survey, candidate metal-poor disk, follow-candidate metal-poor disk, follow-up echelle dataup echelle data

Bulge Matches Thick Disk same massive-star IMFsame massive-star IMF

Gonzales et al 2011

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

Smith et al 2003

Gilmore & Wyse 1991

Local disk

Hiatus then burst Pompeia et al 2008

LMC: solid

dSphs vs. MWG abundances: SFHdSphs vs. MWG abundances: SFH(from A. Koch, 2009 + updates)(from A. Koch, 2009 + updates)

Shetrone et al. (2001, 2003): 5 dSphsShetrone et al. (2001, 2003): 5 dSphs

Letarte (2006): Fornax

Sadakane et al. (2004): Ursa Minor

Koch et al. (2006, 2007): CarinaMonaco et al. (2005): Sagittarius

Koch et al. (2008): HerculesShetrone et al. (2008): Leo II

Aoki et al. (2009): SextansFrebel et al. (2009): Coma Ber, Ursa Major

Hill et al. (in prep): Sculptor

Boo I

◊

Leo IV

Scl

Norris et al 10 BooI

Gilmore et al;

Frebel et al 10 Scl Simon et al 10 Leo IV

Same ‘plateau’ in [Same ‘plateau’ in [αα/Fe] in all systems at /Fe] in all systems at

lowest metallicitieslowest metallicities Type II enrichment only: massive-star IMF Type II enrichment only: massive-star IMF

invariant, and apparently well-sampled/mixedinvariant, and apparently well-sampled/mixed Stellar halo could form from any system(s) in Stellar halo could form from any system(s) in

which star-formation is short-lived, and inefficient which star-formation is short-lived, and inefficient

so that mean metallicity kept lowso that mean metallicity kept low Star clusters, galaxies, transient structures…Star clusters, galaxies, transient structures…

Complementary, independent age information Complementary, independent age information

that bulk of halo stars are OLD further constrains that bulk of halo stars are OLD further constrains

progenitors progenitors (e.g. Unavane, Wyse & Gilmore 1996)(e.g. Unavane, Wyse & Gilmore 1996)

M92 M15

Main sequence luminosity functions of UMi dSph and of globular clusters are indistinguishable.

Wyse et al 2002

HST star counts

0.3M

Star Counts: Invariant Low-Mass IMFStar Counts: Invariant Low-Mass IMF

UMi dSph stars have narrow range of ages, all old

Low-Mass Stellar MF in Low-Mass Stellar MF in Bulge: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..

Stetson et al 2011

Carina dSph CMD

Very extended, non-monotonic star formation history

Carina dSph – extended, bursty star formation history

Carina data: bursts + inhomogeneous star formation

Koch et al inc RW 2008Massive star IMF invariant

Lemasle et al 2012

Age estimates: younger indeed higher [α/Fe]

A much simpler system: Bootes I ‘ultra-faint’ dwarfA much simpler system: Bootes I ‘ultra-faint’ dwarf

SDSS Discovery CMD (Belokurov et al, inc RW, 2006b)Subaru (Okamoto, PhD, 2010)

M* ~ 4 x 104 M, dist ~ 65 kpc

Norris, RW et al 2010

Dwarf spheroidal galaxies have well-defined peaks, with low metallicity tails: self-enriched, from primordial gas? Then lost most baryons – M/L high.

[Fe/H] distributions and radial dependence

Very luminous globular cluster lacks low-metallicity tail; mostclusters do not self-enrich in Fe;Need enough compact baryons

Segue 1, 7 stars

16 stars

Alpha Abundances:Alpha Abundances: 8 stars in Boo I, VLT UVES8 stars in Boo I, VLT UVES Double-blind analysis (Gilmore et al Double-blind analysis (Gilmore et al

2012)2012)

minimal scatter

Boo-119 is CEMP-no star; open dots are field CEMPCEMP-no star Segue 1-7 has [Mg/Fe] ~ 0.94 (Norris et al 2010)

Carbon-enhanced star in Segue 1 (triangles) and BooI (circles)

No s-process plus high [Mg/Fe]

Norris et al 2010a,b

Including data for Boo I stars from Lai et al 2011

[Fe/H] time ISM mixing scale

Two modes of enrichment? Unmixed, very early, enriched by Unmixed, very early, enriched by

individual supernovae from zero-metal individual supernovae from zero-metal stars?stars?

Extremely well mixed, fully sample Extremely well mixed, fully sample massive-star IMF – minimal scatter in massive-star IMF – minimal scatter in element ratioselement ratios

Boo I probably lost 90% of baryons – Boo I probably lost 90% of baryons – metals?metals?

ConclusionsConclusions

Lack of variations in elemental Lack of variations in elemental abundances probably produced by core-abundances probably produced by core-collapse supernovae argue for invariant collapse supernovae argue for invariant massive-star IMFmassive-star IMF Star counts imply fixed low-mass IMFStar counts imply fixed low-mass IMF

Overall patterns determined by star-Overall patterns determined by star-formation historyformation history

Small scatter implies well-mixed ISMSmall scatter implies well-mixed ISM WHY? And HOW? WHY? And HOW?

Large Scale FlowsLarge Scale Flows Chemical evolution plus global star Chemical evolution plus global star

formation rates argue for gas formation rates argue for gas replenishment replenishment

High velocity clouds existHigh velocity clouds exist Galactic FountainGalactic Fountain Cold Flows from Cosmic WebCold Flows from Cosmic Web Accretion from satellite galaxies Accretion from satellite galaxies

(Magellanic Stream)(Magellanic Stream) Radial migrationRadial migration

BoBoöötes Ites I MMVV ~ -6.3, M ~ -6.3, M* * ~ 4 x 10~ 4 x 1044 M M (Kroupa IMF), distance (Kroupa IMF), distance

of ~ 65kpc, half-light radius ~ 250pc (< dark of ~ 65kpc, half-light radius ~ 250pc (< dark matter scalelength?), central velocity dispersion matter scalelength?), central velocity dispersion ~ 3-6 km/s (?), derived (extrapolated) mass ~ 3-6 km/s (?), derived (extrapolated) mass within half-light radius ~ 10within half-light radius ~ 106-76-7 M M, M/L ~ 10, M/L ~ 1033, ,

mean dark matter density ~ 0.1Mmean dark matter density ~ 0.1M/pc/pc3 3

collapse at redshift z > 10collapse at redshift z > 10

Color-magnitude diagram consistent with old, Color-magnitude diagram consistent with old, metal-poor population, similar to classic halo metal-poor population, similar to classic halo globular clusterglobular 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

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

Koposov, et al (inc RW), 2011b

Previous literature value(already reduced)

Identify 37 members, based on line-of-sight velocity, metallicity and stellar gravity (should be giants, dwarfs will be foreground field halo stars)

Koposov, et al (inc RW), 2011b

Field of StreamsField of Streams (and dots)(and dots)

SDSS data, 19< r< 22, g-r < 0.4 colour-coded SDSS data, 19< r< 22, g-r < 0.4 colour-coded by mag (distance), blue (~10kpc), green, red by mag (distance), blue (~10kpc), green, red (~30kpc)(~30kpc)

Belokurov et al (inc RW, 2006)Belokurov et al (inc RW, 2006)

Outer stellar halo is lumpy: but only ~15% by mass Outer stellar halo is lumpy: but only ~15% by mass (total mass ~ 10(total mass ~ 1099MM) and dominated by Sgr dSph ) and dominated by Sgr dSph streamstream

Segue 1

Boo I

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)

Norris, RW et al 2010 Wide-area spectroscopyRed: Segue 1 Black: Boo I

Geha et al І

Wyse & Gilmore 1992

Salpeter IMF slope: -1.35Scalo: -1.5Matteucci for Bulge: -1.1

Chemical Abundances: Boo I & Chemical Abundances: Boo I & Segue 1Segue 1

Spectroscopic surveys with the 2dF/AASpectroscopic surveys with the 2dF/AAΩΩ fibre-fed MOS; stars selected from SDSS to fibre-fed MOS; stars selected from SDSS to follow discovery CMD: wide-area mapping follow discovery CMD: wide-area mapping unique capability of 2dFunique 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 (3933Iron from calibration of Ca II K line (3933Å, Å, as field halo surveys, as field halo surveys, Beers et al 99Beers et al 99), +/- 0.2 ), +/- 0.2 dex dex (Norris et al 08)(Norris et al 08)

Carbon from synthesis of CH G-band, +/- 0.2 Carbon from synthesis of CH G-band, +/- 0.2 in Boo I and +/- 0.4 in Segue 1 (Norris et al in Boo I and +/- 0.4 in Segue 1 (Norris et al 10)10) Follow-up UVES echelle data, [Fe/H] +/- 0.1dex; [C/Fe] for 1

star

Caveat: Segue 1 in very complex part of Caveat: Segue 1 in very complex part of the Galaxythe Galaxy Very flat (bimodal?) metallicity distribution, Very flat (bimodal?) metallicity distribution,

unlike other dwarf galaxies: contamination?unlike other dwarf galaxies: contamination? Extended structure around itExtended structure around it

Same distance and line-of-sight as Sgr stream, Same distance and line-of-sight as Sgr stream, but different velocity but different velocity (Niederste-Ostholt et al wrong (Niederste-Ostholt et al wrong orbit for Sgr stream)orbit for Sgr stream)

Distance and velocity, line-of-sight match Distance and velocity, line-of-sight match Orphan stream Orphan stream (Newberg et al 2010, Koposov et al inc (Newberg et al 2010, Koposov et al inc RW 2011a)RW 2011a)

What is the `300km/s stream’? What is the `300km/s stream’? Extremely wide-field mapping needed to be Extremely wide-field mapping needed to be

assured of statusassured of status

Segue 1Segue 1 MMVV ~ -1.5, M ~ -1.5, M** ~ 600 M ~ 600 M, distance of , distance of

~25kpc, half-light radius ~ 30pc (?), ~25kpc, half-light radius ~ 30pc (?), velocity dispersion ~ 4 km/s (?), derived velocity dispersion ~ 4 km/s (?), derived mass within half-light radius ~ 3 x mass within half-light radius ~ 3 x 101055MM(?), M/L ~ 2000 (?), again <(?), M/L ~ 2000 (?), again <ρρ>>DMDM ~ ~

0.1 M0.1 M/pc/pc3 3 and high collapse redshiftand 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 populationAgain old, metal-poor populationBelokurov et al. 07; Martin et al 08; Geha et al 09; Walker et al 09;Simon et al 2010