SFH, Abundances, and Kinematics of Dwarf Galaxies in the LG Eline Tolstoy, Vanessa Hill, & Monica Tosi 2009, ARA&A, 47, 371 Yin Jun

Content  Introduction  Detailed SFHs  Stellar kinematics and metallicities  Detailed abundances of resolved stars  Chemical evolution models

1. Introduction  What is a dwarf galaxies? Fainter than M B =-16 (or M V =-17), more extended than GC (e.g., Tammann 1994); Presence of DMH (e.g., Mateo 1998); Types: dSphs, dIs, uFd, BCDs, UCDs… Is there any other physical property that distinguishes a dwarf galaxy from bigger galaxies?

1. Introduction  Es, GCs: Clear distincted;  Early- and late-type dwarfs, BCDs fall along similar relations;  Overlap with larger late type and spheroidals;  uFds: clearly separated but arguably follow the relation;  No clear separation between dwarf galaxies and the larger systems;  dIs, BCDs, and dSphs tend to overlap with each other in this parameter space. Early-type dwarfs are the same as late-type systems that have been stripped of their gas.

1. Introduction  What is a dwarf galaxies? Fainter than M B =-16 (or M V =-17), more extended than GC (e.g., Tammann 1994); Presence of DMH (e.g., Mateo 1998); Types: dSphs, dIs, uFd, BCDs, UCDs… Is there any other physical property that distinguishes a dwarf galaxy from bigger galaxies? NO! These low-metallicity systems show a wealth of variety in their properties, such as luminosity, surface brightness, SFH (both past and present), kinematics, and abundances.

1. Introduction  dwarf galaxies in LG: Individual stars can be resolved and photometered down to the oldest MSTOs.  provides the most accurate SFHs going back to the earliest times. Spectra can be taken of individual RGB stars  chemical content & kinematics of a stellar population. The most accurate studies of resolved stellar populations have been made.  Understand the relation with larger cousins, ISM, diff. types, GCs… (Mateo 2008)

2. Detailed SFH 2.1 Techniques: Synthetic CMD Analysis

 It is based on comparing observed with theoretical CMDs created via Monte-Carlo-based extractions from stellar evolution tracks, or isochrones, for a variety of star-formation laws, IMFs, binary fractions, age- metallicity relations, etc. Photometric errors, incompleteness, and stellar crowding factors also have to be estimated and included in the procedure to fully reproduce an observed CMD.  reproduces all the main features of the observational one: morphology, luminosity, color distribution, and number of stars in specific evolutionary phases.

2. Detailed SFH 2.1 Techniques: Synthetic CMD Analysis  Same: n * =50,000 Salpeter IMF  Different: Z SFH  Key: Z degeneracy interpolating

2. Detailed SFH 2.1 Techniques: Synthetic CMD Analysis  Reliability: three different people; different assumptions, modeling procedures, and even stellar evolution models; consistent results within their uncertainties.

2. Detailed SFH 2.2 Observations: Dwarf Galaxies in the Local Group Holmberg limit faintest main feature visible in the CMD explicitly not detected HB and/or oldest MSTOs ancient population: RR Lyr variable stars CMD was not deep enough to determine Individual RGB stars LR: R<10,000 HR: R>18,000

2. Detailed SFH 2.2 Observations: Dwarf Galaxies in the Local Group

(Mateo 2008) GC closest dwarfs dwarfs

2. Detailed SFH 2.2 Observations: Dwarf Galaxies in the Local Group indication of the possibility of morphological transformation (Mateo 2008) Spirals dSph remote dSph transition dIrr

Early-type dwarf galaxies  typically associated with large galaxies  closest to us, the majority at distances < 130 kpc  look very much like the old extended stellar pop., which underlie most late-type systems. major difference: lack gas and recent SF  compact dEs: M32-like galaxies, rare low luminosity elliptical, not because of tidal pruning  diffuse dSphs: NGC205-like galaxies, common  UCDs: May be like ωCen Tidally stripped nucleus of a compact system 2. Detailed SFH 2.2 Observations: Dwarf Galaxies in the Local Group

completely distinct episodes of star formation

Late-type dwarf galaxies  probe metal-poor SF, both young and old;  retain HI gas and are typically forming stars at the present time with a variety of rates;  numerous and often fairly luminous class within LG;  typically at distances >400 kpc (except SMC) Two examples:  Leo A: deepest and most accurate ever make for dI  LGS 3: transition-type, contains HI gas, no very young stars (no HII regions, no supergiant) 2. Detailed SFH 2.2 Observations: Dwarf Galaxies in the Local Group

dSph transition-type dI 90% 775 kpc620 kpc800 kpc

2. Detailed SFH 2.2 Observations: Dwarf Galaxies in the Local Group Ultrafaint dwarf galaxies  M V -r 1/2, offset; M V - μ V, extension of dSph;  However, exist in a region where both the extension of classical dwarfs and the GCs sequences may lie  Absolute magnitude: -8<M V <-1.5  most of uFds have been found in the immediate vicinity of the MW most distant: Leo T (410 kpc) and CVn I (218 kpc) typical distances: 23 kpc (Seg I) to 160 kpc (Leo IV, CVn II)  Synthetic CMD method: difficult to distinguish stars  look for distinctive stellar populations: blue HB (BHB) stars or RR Lyr variable stars

2. Detailed SFH 2.2 Observations: Dwarf Galaxies in the Local Group SMC  closest late-type dwarf (~ 60 kpc);  high gas content, low Z (~0.004), low mass (1~5×10 9 M ⊙ )  hosts several hundred star clusters covering ages from 11 Gyrs (NGC 121) to a few Myrs (NGC 346 and NGC 602)  less studied than might be expected Few ground-based and HST-based studies on small regions Extensive surveys of whole SMC are planned  stars older than 8 Gyears do not dominate the SMC population, the population bulk seems to peak at ages somewhat younger than 6–9 Gyears essentially everywhere in the SMC main body. the cluster has formed most of its stars around 2.5 Myears ago, whereas the surrounding field has formed stars continuously since the earliest epochs.

2. Detailed SFH 2.3 Beyond the Local Group  actively star-forming BCDs (e.g., I Zw 18, NGC 1705).  The further the distance, the worse the crowding conditions and the shorter the look-back time reachable  there is no evidence that any of these systems is younger than the look-back time.  SFH studies both in the Local Group and beyond Vast majority of dwarfs have fairly moderate SF activity extensive Hα study of 94 late-type galaxies:  the typical SFR of irregular galaxies is 10 −3 M ⊙ yr −1 kpc −2  BCDs is generally higher but not by much.  the star-formation regions are not intrinsically different in the various galaxy types, but they crowd more closely together in the centers of BCDs. BCD least active BCD dI most active LSB

3. Stellar Kinematics and Z 3.1 Early-Type Dwarfs  Close, RGB stars: trace history  Expected σ< 2 km s -1, observed σ~ 8-15 km s -1 Contain significant amount of DM Or do not understand gravity in these regimes  It was found that RGB stars of a different metallicity range (and, hence, presumably age range) in dSphs can have noticeably different kinematic properties  MDF Use CaII triplet metallicity indicator, fail at low Z ([Fe/H]<-2.5) Compare dSph with Galactic halo Challenge to models where where all of the Galactic halo builds up from the early merging of dwarf galaxies

3. Stellar Kinematics and Z 3.2 Late-Type Dwarfs  Far away, HI gas & HII region: present time  Kinematics: HI gas influenced by on-going star-formation processes. HI velocity dispersion is almost always ∼ 10 km s −1 in any system, from the smallest dIs to the largest spiral galaxies difficult to compare the kinematic properties of dIs and dSphs. LGS 3 & Leo T: HI & stars, no sign of rotaion  Metallicity: spectroscopy of massive stars or HII region only a few million years old in dSphs, typically measured for stars older than ∼ 1 Gyr difficult to accurately compare early- and late-type dwarfs in all dIs, the HII regions in a single galaxy appear to have identical [O/H] abundances

3. Stellar Kinematics and Z 3.3 Ultrafaint Dwarfs  The stellar kinematics and metallicities play an important role quantify the degree of disruption faint galaxies or some kind of diffuse GCs Difficult: embedded in the foreground of our Galaxy, both in position and in velocity  Kinematics more dark matter–dominated, M/L ∼ 140–1700. did not correct for tidal effects  Metallicities lower than in most GCs ([Fe/H] ≤−2), and with a larger scatter also lower than in other more luminous dwarf galaxies found C-rich metal-poor star, similar to those in the MW halo  Brighter uFds (M V <-5): low-mass tail to dSphs and dI/transition  Fainter uFds: lie in gap between GCs and dwarfs If the large DM masses are correct  an extension of the galaxy class

4. Detailed abund. of resolved stars 4.1 Dwarf Spheroidal Galaxies Alpha elements  Mearsured in RGB spectra: O, Mg, Si, Ca, Ti O, Mg: produced during the hydrostatic He burning in massive stars Si, Ca, Ti: produced during the SN II explosion  “knee” the time SNe Ia start to contribute to the chemical evolution This is between 10 8 and 10 9 years after the first SF episode. Stop SF 10 Gyrs ago Age<1Gyr steady SF 2~10 Gyrs 3 bursts

4. Detailed abund. of resolved stars 4.1 Dwarf Spheroidal Galaxies Alpha elements  The position of “knee” expected to be different for different dSphs. correlates with the total luminosity and the mean metallicity. higher [Fe/H]: efficiently produces and retains metals lower [Fe/H] : either loses significant metals in a galactic wind, or simply does not have a very high SFR. metal-poor side of the knee: indistinguishable from those in MW halo metal-rich side of the knee: decrease of [α/Fe] with increasing metallicity  Lower [α/Fe] in dSph than in MW disk or halo  sudden decrease of SF, result of galactic winds or tidal stripping.

4. Detailed abund. of resolved stars 4.1 Dwarf Spheroidal Galaxies Sodium and Nickel  Na mostly produced in massive stars (during hydrostatic burning) with a metallicity-dependent yield. [Fe/H]<-1: no apparent difference btwn dSph and MW halo stars [Fe/H] >-1: dSph produce too little Na  Ni also largely produced in SNe Ia  [Na/Fe]-[Ni/Fe] correlation tentatively explained as the common sensitivity of both elements to neutron-excesses in supernovae. can be modified by SN Ia nucleosynthesis, especially in high Z low [α/Fe] populations of dwarfs Too little Na in dSphs No apparent difference ?

4. Detailed abund. of resolved stars 4.1 Dwarf Spheroidal Galaxies Neutron-capture elements  Nuclei heavier than Z ∼ 30 s -process: occur in low- to intermediate-mass (1–4 M ⊙ ) AGB stars, delay time is ∼ 100–300 Myrs r-process: massive-star, the most plausible candidates are SNe II, delay time is very little.  Y, Ba, La: either the s- or the r- process  Eu: r- process only  In the MW, Ba and Y are dominated by r-process for [Fe/H]< −2.0 s-process for [Fe/H]> −2.0 little difference s-process dominates r-process dominates 1. dwarfs enriched faster than the halo at the earliest times 2. the site for the r -process is less common (or less efficient) in dSphs r-process only

4. Detailed abund. of resolved stars 4.1 Dwarf Spheroidal Galaxies s-process dominates Ba: r-process dominates Eu: r-process only [α/Fe] knee

4. Detailed abund. of resolved stars 4.2 Ultrafaint Dwarf Galaxies The overall similarity between all the most metal-poor stars for element ratios up to the iron-peak can be taken as an indication that SF and metal-enrichment, even at the earliest times, and even in the smallest systems, has proceeded in a similar manner. discrepancy

4. Detailed abund. of resolved stars 4.3 Dwarf Irregulars  large distance: HII region and supergiant look-back time: at most a few tens of Myrs present-day metallicity are all more metal poor than MW disk young population (L-Z relation) ~7.3 < 12+log(O/H)< ~8.1 little dispersion within galaxy, no spatial gradient  suggest a very efficient mix of metals across the galaxy despite the clumpiness of ISM and ongoing SF.  The shear is very low, mixing occurs in the gaseous hot phase typically only light elements (e.g., He, N, O, Ne, S, Ar), no iron (nor any other element that would trace SNIa) Lower [α/Fe] than in larger systems (e.g. MW, LMC)  low SFR and/or metal losses through winds dIs actually prolong the trends of dSph galaxies, dSphs are entirely consistent with dIs that lost their gas at a late stage of their evolution. The Fnx dSph and the SMC, which are both dominated by intermediate-age populations, are also quite similar in their chemical enrichment, except that Fnx ran out of gas (or lost its gas) and stopped star formation about 10 8 years ago.

5. Chemical evolution models 5.1 Explaining Low Metallicity  variations in the IMF; steeper IMF slopes and/or mass range cut-offs have been proposed to reduce the chemical enrichment from massive stars;  accretion of metal-free, or very metal-poor, gas to dilute the enrichment of the galaxy;  metal-rich gas outflows, such as galactic winds, triggered by supernova explosions in systems with shallow potential wells, or gas stripping due to interactions with other galaxies or with the IGM to efficiently remove the metal-enriched gas from the system.

5. Chemical evolution models 5.2 Galactic Winds  have been predicted by hydrodynamical simulations to be able to remove a large fraction of the elements synthesized by SNe II as well as a fraction of the galaxy’s interstellar medium;  there is increasing observational evidence for starburst- driven metal-enriched outflows;  Explain the low metallicity  Naturally explain the L-Z relation  Explain the structural similarities observed by Kormendy (1985)  The influence of tidal effects is considered to play an important role, but hard to verify.

5. Chemical evolution models 5.3 Modeling Individual Systems Standard chemical evolution models  Take into account global parameters, follow the evolution of individual element  Simplistic assumption on dynamics  successful in predicting large-scale, long-term phenomena  simplistic treatment of stellar and supernova feedbacks and of gas motions, is an obvious drawback. Chemodynamical models  analyze in detail the heating and cooling processes and put important constraints on the onset and fate of galactic winds, stripping, and ram pressure  dynamics processes in great detail  successful in predicting small-scale, short-term phenomena  cannot follow galactic-scale evolution over more than a gigayear Carigi, Hernandez & Gilmore (2002) Romano, Tosi & Matteucci (2006) Carigi, Colin & Peimbert (2006), Lanfranchi, Matteucci & Cescutti (2008) Recchi et al. (2004, 2006) Fenner et al. (2006) Marcolini et al. (2006, 2008) The challenge in the next few years is to improve both types of approaches and get a more realistic insight into how stars and gas evolve, chemically and dynamically, in their host galaxies.

Thank You