Stellar Feedback EffectsStellar Feedback Effectson Galaxy Formationon Galaxy Formation

Filippo Sigward

Università di FirenzeDipartimento di Astronomia e Scienza dello Spazio

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

Andrea Ferrara, SISSA / ISAS

Evan Scannapieco, KITP, SB

Why Feedback ?Why Feedback ?

Ingredients for Galaxy formation and evolution:

• Evolution of dark halos• Cooling and star formation • Chemical enrichment• Stellar populations

Model outputs

Comparison with observations

• Feedback

The “cooling catastrophe”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

Early preheatingEarly preheating

Benson & Madau 2003

Increased gas pressure by winds from pregalactic starburst & energy deposited by accreting BH.

Global early energy input: “preheating”

LF

ObservedGood agreement

in the faint-end slope

Unable to explain the cut-off at bright magnitudes

Additional feedback processes to suppress dwarf galaxies: SN-driven shocks from nearby galaxies

Previous Analytical StudiesPrevious Analytical Studies

Ctn

kTcool

e

s 2

3

• Mechanical evaporationMechanical evaporation:

Ts > Tvir

– Cooling:

• Baryonic strippingBaryonic stripping:

f Ms vs Mb ve

(Scannapieco, Ferrara & Broadhurst 2000)

CDM

CDM

Numerical simulationsNumerical simulations

• Pre-virialized case: Bertschinger 1985 (analytical and semi-analytical

solutions)

• Virialized case: Navarro, Frenk & White 1997 (cosmological simulations)

Initial conditionsInitial conditions

- Shock:

1 SN occurs every 100 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

Pre-virialized casePre-virialized case

distance [kpc]

b[

g cm

–3]

distance [kpc]

v [c

m s

–1]b r –2.25

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

37

9

10 18

M

z

hMM

Pre-virialized case

kpc735.taR

Simulation parameters:Initial density [g cm–3]

x 138 pc

20.7 kpc

Final mapsFinal mapsPre-virialized case

Density [g cm–3] Temperature [K]

t = 133 Myr

Rta Rta

22000

02

exp0 22

cb

eevir

pbb r

kT

mr

vv

- Dark Matter profile (NFW):

- Baryonic profile:

Virialized caseVirialized case

c

vir

cc

c

R

rx

cxcxr

84

1 2

.

characteristic overdensity

b [

g cm

–3]

distance [kpc]

14

9

10 18

M

z

hMM

kpc751.virR

K12490virT

138 skmvire Rv

Simulation parameters:Initial density [g cm–3]

Virialized case

x 43 pc

6.5 kpc

Density maps: evolutionDensity maps: evolutionVirialized

case

6.5 kpc

time: 0 - 58.2 Myr

Final mapsFinal mapsVirialized case

t = 58.2 Myr

Density [g cm–3] Temperature [K]

RvirRvir

Amount of Gas RemovedAmount of Gas Removed

Mb (T > Tvir) / Mb ~ 5.0%

Mb (v ve ) / Mb ~ 69.9%

tf = 133 Myr

Pre-virialized case total

v ve

T > Tvir

t [Myr]

Mbou

t (t)

/ Mb(

t)

Mbou

t (t)

/ Mb(

t)

t [Myr]

T > Tvir

v ve

Mb (T > Tvir) / Mb ~ 0.9%

Mb (v ve ) / Mb ~ 0.7%

tf = 58.2 Myr

Virialized case total

Amount of Gas RemovedAmount of Gas Removed

ConclusionsConclusions

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.

1. Strong suppression of dwarf galaxy formation by shocks from nearby galaxies can occur in the collapse stage immediately after the turn-around.