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Stellar Feedback Effects on Galaxy Formation

Japan – Italy Joint Seminar “Formation of the First Generation of Galaxies: Strategy for the Observational Corroboration of Physical Scenarios” December 2 – 5, 2003 – Niigata University, Japan. Stellar Feedback Effects on Galaxy Formation. Filippo Sigward Università di Firenze

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Stellar Feedback Effects on Galaxy Formation

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  1. Japan – Italy Joint Seminar “Formation of the First Generation of Galaxies: Strategy for the Observational Corroboration of Physical Scenarios” December 2 – 5, 2003 – Niigata University, Japan Stellar Feedback Effectson Galaxy Formation Filippo Sigward Università di Firenze Dipartimento di Astronomia e Scienza dello Spazio Andrea Ferrara, SISSA / ISAS Evan Scannapieco, KITP, SB

  2. Why Feedback ? Ingredients for Galaxy formation and evolution: • Evolution of dark halos • Cooling and star formation • Chemical enrichment • Stellar populations • Feedback  Comparison with observations Model outputs

  3. The “cooling catastrophe” In the absence of any contrasting effect, much of the gas is expected to sink into small halos at early epochs  Strong feedback is invocated to avoid too many baryons turning into stars at primeval ages

  4. Benson & Madau 2003 Early preheating Increased gas pressure by winds from pregalactic starburst & energy deposited by accreting BH. Global early energy input: “preheating” LF Unable to explain the cut-off at bright magnitudes Observed Good agreement in the faint-end slope Additional feedback processes to suppress dwarf galaxies: SN-driven shocks from nearby galaxies

  5. Previous Analytical Studies CDM •Mechanical evaporation: Ts> Tvir – Cooling: CDM • Baryonic stripping: f Msvs Mbve (Scannapieco, Ferrara & Broadhurst 2000)

  6. Numerical simulations • Pre-virialized case:Bertschinger 1985 • (analytical and semi-analytical solutions) • Virialized case: Navarro, Frenk & White 1997 • (cosmological simulations)

  7. Initial conditions • Shock: • 1 SN occurs every100 Mof baryons that form stars • sf = 0.1 • Etot / SN = 2 1051 erg • Outflows initialization: thin shell approximation • Rs = mean distance between the halos • plane wave (Rs Rvir,ta) • IGM:igm homogeneous, T = 104 K

  8. v [cm s–1] b[g cm–3] distance [kpc] distance [kpc] Pre-virialized case Similarity solutions for infall and accretion onto an overdense perturbation (Bertschinger 1985). Rtat8/9 M(r < Rta)  t2/3 Particles come to rest after the shock b r –2.25

  9. Initial density [g cm–3] Simulation parameters: Pre-virialized case x 138 pc 20.7 kpc

  10. Final maps Pre-virialized case t = 133 Myr Rta Rta Density [g cm–3] Temperature [K]

  11. b [g cm–3] characteristic overdensity distance [kpc] Virialized case - Dark Matter profile (NFW): - Baryonic profile:

  12. Initial density [g cm–3] Simulation parameters: Virialized case x 43 pc 6.5 kpc

  13. Density maps: evolution Virialized case 6.5 kpc time:0 - 58.2Myr

  14. Final maps Virialized case t = 58.2 Myr Rvir Rvir Density [g cm–3] Temperature [K]

  15. total v ve Mbout(t) / Mb(t) T > Tvir t [Myr] Amount of Gas Removed Pre-virialized case Mb (T > Tvir) / Mb~ 5.0% tf = 133 Myr Mb (v ve) / Mb~ 69.9%

  16. T > Tvir Mbout(t) / Mb(t) v ve t [Myr] Amount of Gas Removed Virialized case total Mb (T > Tvir) / Mb~ 0.9% tf = 58.2 Myr Mb (v ve) / Mb~ 0.7%

  17. Conclusions • Strong suppression of dwarf galaxy formation by shocks from nearby galaxies can occur in the collapse stage immediately after the turn-around. 2. Such feedback is much less efficient (a few % mass loss) if the system is already virialized. 3. Gas is predominantly removed via baryonic stripping; mechanical evaporation is not efficient due to rapid cooling of the halo gas.

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