Chemical Evolution of Galaxies. Francesca Matteucci, Trieste University Russbach, 15 March 2011

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1 Chemical Evolution of Galaxies Francesca Matteucci, Trieste University Russbach, 15 March 2011

2 Outline of the talk Basic ingredients for building chemical evolution models Examples of chemical abundance patterns which can constrain galaxy formation and evolution: The Milky Way The dwarf spheroidals of the Local Group Elliptical galaxies

3 Chemical Evolution Chemical abundances tell us about the nucleosynthesis as well as the formation and evolution of galaxies Light elements (H, D, He, Li) were synthesized during the Big Bang All the elements with A>12 were formed inside stars Stars produce new elements and then restore them into the ISM. This process is chemical evolution

4 Basic Ingredients of Chemical Evolution Initial conditions (open, closed-box; initial chemical composition) The stellar birthrate function: SFRxIMF The stellar yields (i.e. the mass restored into the ISM by a star of a given mass in the form of a given chemical element Gas flows: infall, outflow, inflow

5 The Star Formation History The main driver of chemical evolution in a galaxy is its star formation history The star formation history depends on the SFR (star formation rate) and IMF (initial mass function) The IMF describes the distribution of stellar masses at birth, the SFR tells how much gas goes into stars per unit time

6 The Initial Mass Function The IMF has been derived for the solar vicinity. It is a power law:

7 The parametrization of the SFR The most common parametrization is the Schmidt (1959) law, where the SFR is proportional to some power (k=2) of the gas density Kennicutt (1998) suggested k=1.5 from studying star forming galaxies, but also a law depending on the rotation angular speed of gas and the first power of the gas density

8 The Stellar Yields Low and intermediate mass stars (0.8-8 Msun): produce He, N, C and heavy s-process elements. They die as C-O WDs, when single, and can die as Type Ia SNe when binaries Massive stars (M>8-10 Msun, core-collapse SNe): they produce alpha-elements (O, Mg..), some Fe, light s-process elements and r- process elements Type Ia SNe produce mainly Fe ( M_sun per SN)

9 Supernova Progenitors Core-collapse SNe originate from massive stars, they can be of Type II,Ib and Ic. They leave neutron stars or black holes as remnants. Type Ib/c SNe are connected to GRBs Type I a SNe are thought to originate from C-O white dwarfs (WD) in binary systems. They do not leave any remnant

10 Type Ia Supernovae Double-Degenerate scenario (Iben & Tutukov, 1984): two C-O WDs merge after loosing angular momentum due to gravitational wave radiation When the two WDs merge and the Chandrasekhar mass is reached, then an explosion occurs. Minimum time for the first system to explode is 35Myr (lifetime of 8Msun)+ gravitational time delay (1Myr, Greggio 2005)

11 Type Ia Supernovae Single-degenerate scenario (e.g. Whelan & Iben 1974; Hachisu et al. 1996; Han &Podsiadlowski 2004): a binary system with a C-O white dwarf plus a MS star. When the star becomes RG it starts accreting mass onto the WD When the WD reaches the Chandrasekhar mass, it explodes by C-deflagration as Type Ia supernova. Minimum time for the first system to explode is 35 Myr!

12 Gas flows Gas can be accreted or lost from a galaxy The most common parametrization of the accretion rate is an exponential law

13 Gas Flows Outflows and galactic winds are seen in galaxies The most common parametrization is Wind=CxSFR The wind rate is proportional to the star formation rate (Martin, 2000,2004)

14 Basic Equations All these ingredients should be included in equations describing the evolution of the abundances of single chemical elements in the gas These equations should be solved numerically if one wants to take into account in details stellar lifetimes

15 Example of Equations

16 The Formation of the MW from Astroarchaelogy The two-infall model of Chiappini, FM & Gratton (1997) predicts two main episodes of gas accretion During the first one the halo, bulge and all or part of thick disk formed, the second gave rise to the thin disk This model can reproduce most of the observational data

17 Recipes for the Two-Infall Model SFR- Kennicutt s law with a dependence on the surface gas density (exponent k=1.5) plus a dependence on the total surface mass density (feedback). Threshold of 7 Msun/pc^2 in thin disk Type Ia SN progenitors from the SD model (Greggio & Renzini 1983; FM & Recchi 2001). This scenario and the DD are the best to reproduce chemical properties (FM &al. 2006;2009) IMF, Scalo (1986) normalized over a mass range of Msun

18 Recipes for the Two-Infall Model Exponential infall law with different timescales for inner halo and thick disk(1-2 Gyr) and thin disk (inside-out formation with 7 Gyr at the S.N.) The efficiency of SF is 2 Gyr^(-1) in the halothick disk and 1Gyr^(-1) in the thin disk The thin disk is divided in several independent rings 2 Kpc wide Radial gas flows, if present, should have a low speed. If included in the model, they contribute to steepen the abundance gradient along the thin disk (e.g. Schonrich & Binney 2009; Spitoni & FM 2011)

19 Results for the MW (Francois, FM & al. 2004):solar vicinity

20 Effect of different yields Predicted and observed [X/Fe] for two different sets of yields Dashed lines:yields from Woosley & Weaver (1995) for massive stars and from van den Hoeck & Groenewegen (1997) for low and intermediate mass stars Continuous lines: yields from Karakas (2010) for low and intermediate mass stars and from Geneva group for He, C, N,O and from Kobayashi & al. ( 2006) for heavt elements Models from Romano & al Best yields: continuous line

21 Romano & al Best yields: dashed line

22 Heavy elements Predicted and observed [Ba/Fe] in the Milky Way The red line is the prediction of Cescutti et al. (2006) with Ba originating from Msun The rest of Ba is the s-component and comes from 1-3 Msun (Busso & al. (2001) Data from Francois et al. (2007)

23 The G-dwarf Metallicity Distribution G-dwarf metallicity distribution compared with predictions of the two-infall model (Kotoneva et al. 2003) The assumed timescale for disk formation is a very important parameter

24 G-dwarf Metallicity Distribution: the effect of the IMF More recent data CGS (Nordstrom et al. 2004; Holmberg et al. 2007) indicated by black line Green line: prediction with Scalo IMF. Red line: prediction with Kroupa et al. (1993) IMF

25 The time-delay model and the star formation rate in galaxies Predicted [alpha/fe] ratios for different SFR histories (a more modern version of FM & Brocato 1990) A strong starburst (dashed line), a SFR like in the solar vicinity (dotted) and a slow SFR (continuous) like in Magellanic Irregulars or Dwarf Spheroidals

26 Thin and Thick Disk More recent data (Ramirez &al. 2007; Bensby & al. 2005; Melendez & al. 2008) show two distinct parallel sequences for the [O/Fe] ([Fe/O]) ratio Thick disk stars show [O/Fe] larger than thin disk stars (magenta) Chiappini (2008) made the thick disk to form with SFR 10 times the one of the S.N.

27 The Galactic Bulge Model (black, Cescutti & FM 2010): fast Bulge formation (0.3 Gyr) and IMF flatter than in S.N. Predicts large Mg to Fe for a large Fe interval Turning point at larger than solar Fe (effect of the time-delay model) Data from giants and dwarfs. Bensby & al. 2010; Alves-Brito & al. 2010

28 The Galactic Bulge Models with different IMFs The best IMF for the Bulge is flatter than in the S.N. (black line) and flatter than Salpeter (blue line) (Ballero & al. 07) Best IMF: x=0.95 for M> 1 solar mass and x=0.33 below

29 Bulge vs. Thick Disk (Fulbright & al. 2007)

30 Bulge vs.thick Disk Bulge stars have higher [O/Fe] and especially [Mg/Fe] than the thick disk stars (Fulbright et al. 07; Lecurer et al. 07) Bulge stars (red) have the same [O/Fe] of thick disk stars (blue) (Melendez et al. 2008)

31 [O/Mg] in the Bulge The steep decrease of the [O/Mg] vs. [Mg/H] in the Bulge can be reproduced by yields with mass loss by massive stars (Maeder 1992) McWilliam, FM & al. (2008)

32 [C/O] in the Bulge Model without mass loss (WW95) does not fit the data! Model with mass loss from massive stars from Maeder (1992) is the best Data and models from Cescutti & al. 09

33 Abundance Gradients Predicted and observed abundance gradients from Chiappini, FM & Romano (2001). Disk forms inside-out Data from HII regions, PNe and B stars, red dot is the Sun The gradients slightly steepen with time (from blue to red). No clear indication from data

34 Abundance Gradients More recent data on gradients include abundances from Open Clusters, Cepheids and O, B stars and show flat gradients To reproduce the gradients inside-out formation of the disk plus threshold in gas density for SF (Colavitti & al. 2008)

35 Time-scales from Galactic Astroarchaelogy The inner stellar Halo must have formed on a timescale of 1.5 Gyr whereas the outer Halo must have formed on a longer timescale (insideout Halo formation) The Local disk must have assembled by gas accretion on a time scale from 6-8 Gyr and the timescale increases with galactocentric radius The thick disk must have formed more quickly than the thin disk (2 Gyr?) but smoothly The Bulge must have formed on a time no longer than Gyr with flatter IMF

36 Summary Time-delay model well explains the disk and Bulge abundance ratios To reproduce the abundance gradients an inside-out disk formation and a threshold in the SF are required There is not yet full agreement on whether the [alpha/fe] ratios in the Bulge are higher or the same as in the thick disk If they are the same, the IMF in Bulge and thick disk should be the same

37 [Alpha/Fe] ratios in Dwarf Spheroidals

38 Chemical evolution of dsphs: standard model Lanfranchi & FM (2003, 2004) proposed a model which assumes a SFH as derived by the CMDs. Initial baryonic masses 5x10^(8)Msun SN feedback induces a strong outflow. DM ten times LM but diffuse (M/L today of the order of 100) SFR less efficient than in the MW and going on for 8 Gyr

39 Specific Models for dsphs: Carina The best model for Carina by Lanfranchi, FM & Cescutti ( 2006) compared with data from Koch & al.(2008) Four bursts of SF are considered, as suggested by the C-M diagram

40 Specific Models for dsphs: Carina The stellar metallicity distribution predicted for Carina and observed by Koch & al. (2005) SF history from Rizzi et al. 03. Four bursts of 2 Gyr, SF efficiency 0.15 Gyr^(-1) < 1-2 Gyr^(-1) (S.N.), wind=7xsfr Salpeter IMF

41 Specific Models for dsphs:sculptor Lanfranchi, FM & Cescutti (2006) predicted the evolution of alpha and heavy elements in dsphs (Ba, Eu, Sr, La). Prescriptions for s- and r- process elements from Cescutti & al. (2006). One extended episode of SF, strong wind (13xSFR), eff.sf 0.2Gyr^(-1) Good agreement between data of Sculptor from Geisler & al. (2007) and model predictions

42 Sculptor: data from Geisler & al. (2007), model Lanfranchi &al.06

43 Sculptor: recent data from Kirby et al Data for Sculptor (Kirby et al. 2009, Shetrone et al. 2003, Geisler et al. 2005)) Models from Lanfranchi & Matteucci (2009) in red

44 Summary The chemical evolution of dsphs in general can be reproduced by assuming an extended period of SF, a poorly efficient SFR and a quite efficient galactic wind from SN feedback The [X/Fe] ratios of alpha and heavy elements suggest different histories of SF between the MW and dsphs Same [alpha/fe] ratios at [Fe/H]<-3.0 dex is not necessarily a proof that the MW halo has formed by accretion of dsphs

45 Elliptical Galaxies Elliptical galaxies are assumed to have formed stars very quickly with intense SF, such as the true bulges Star formation is soon stopped in E galaxies by the occurrence of galactic winds (feedback) which devoid them of gas DM halos 10 times more massive than the stellar mass are considered. Galactic winds develop much before 1 Gyr (Pipino & FM, 2004)

46 Star Formation Histories FM (1994) and Pipino & FM (2004) suggested that SF efficiency increases with galactic mass Thus galactic winds develop first in more massive galaxies (inverse wind scenario) Feedback from SNe is taken into account We can call it downsizing in SF

47 SN Rates in Ellipticals Predicted Type II (dotted) and Type Ia (continuous) SN rates in an elliptical galaxy of 10^(11) Msun luminous mass (Pipino & FM 2004) The galactic wind occurs at 0.4 Gyr, then SF stops. Most of stars will have high [alpha/fe] ratios

48 SF Histories in different environments: Thomas (2006)

49 [alpha/fe] ratios in ellipticals The [alpha/fe] ratios increase with galactic mass in ellipticals Downsizing in SF produces naturally this results. A consequence of the time-delay model Figure adapted from Thomas et al. (2002)

50 Mass-metallicity relation Mass-metallicity relation (real [Mg/H] and [Fe/H]) predicted for ellipticals (right) M-Z relation by means of indices:<fe> and Mg2. The effects of using different index calibrations are shown Predictions from Pipino & FM (2004). Data from Carollo & al. (1993), Trager (1998),Kuntschner & al. (2001) and others

51 Metallicity Gradients in Ellipticals Galactic winds develop first in the external regions of ellipticals (Martinelli et al 1998; Pipino, FM & al. 2006) After the wind, star formation stops The outer regions should be slightly older than the inner ones In principle, [alpha/fe] ratios should increase toward the outer regions

52 Metallicity Gradients in Ellipticals Hydrodynamical simulations in 1-D (Pipino, D Ercole & FM 2008) confirm the outside-in formation Gradients in abundance ratios ([O/Fe]) can be positive, null or negative, depending on the play between radial flows and duration of SF Observations confirm a variety of slopes

53 Summary In order to have a high [Mg/Fe] in the dominant stellar population, ellipticals must have formed stars very quickly (less than 1 Gyr) Galactic winds halt the SF and develop outsidein. Negative abundance gradients but various abundance ratio gradients can develop In order to have [Mg/Fe] increasing with M, the star formation should have lasted for a shorter period in massive than in small ellipticals Down-sizing in SF preserves the M-Z relation

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