primordial star formation: constraints on the imf from protostellar feedback jonathan c. tan eth...
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Primordial Star Formation: Constraints on the IMF from Protostellar Feedback
Jonathan C. TanETH Zurich
Christopher F. McKeeUC Berkeley
Eric G. BlackmanUniversity of Rochester
Image from Scientific American and V. Bromm
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Reionization
Metal Enrichment
Illumination
Progenitors of SN and GRBs?
Influence on Quasars, Globular Clusters, Galaxies?
“Simple” problem: initial conditions, chemistry,no feedback from other stars, weak B-fields(?)
First Stars and the Dawn of Complexity
depend onstellar mass (IMF)
CMB polarization (WMAP Kogut et al. 03)H 21cm (LOFAR Morales & Hewitt 03)
Z of halo stars (Christlieb et al. 03)
NIR bkg. intensity (Santos et al. 02NIR bkg. fluctuations (Kashlinsky et al. 04)
Observations
(Tumlinson, Venkatesan
& Shull 2004)
JWST (Weinmann & Lilly 2005)SWIFT (Bromm & Loeb 2002)
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Hydrogen Ionizing Luminositiesalong the Primordial Main Sequence
Tumlinson & Shull 00; Bromm et al. 01;Ciardi et al. 01; Schaerer 02
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Initial mass determines nucleosynthetic yields and final remnant
Heger & Woosley 2002
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Near Infra-RedBackground
Potential signature of firststars in the NIR EBL (Salvaterra & Ferrara 02)
But zodiacal subtraction isuncertain.
EBL
galaxies
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Overview of Structure Formation
1. Recombination z ≈1200, start of “dark ages”2. Thermal equilibrium matter-CMB until z ≈160
: independent of ze.g. globular clusters
3. Thermal decoupling,
Madau (2002)
4. “First Light”5. Reionization,
e.g. galaxies
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Numerical Simulations: Methods
Abel, Bryan, Norman (2002):
Bromm, Coppi, Larson 1999; Abel, Bryan, Norman 2000,2002
Eulerian AMR; Riemann solvernon-eq. chemistry of 9 speciesoptically thin radiative losses: line cooling Compton heating/cooling
tot=1, =0 , b=0.06,z=100scale invariant power spectrum(128kpc)3 comovingIdentify region of 1st DM halowith ~106Msun, then focus herewith DM particle mass =1Msun
Grid refinement to resolve: Jeans mass, density contrasts, and cooling timescales
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Numerical Simulations: Results
Abel, Bryan, Norman (2002)
1. Form pre-galactic halo ~105-6Msun
2. Form quasi-hydrostatic gas coreinside halo: M≈4000Msun,r ≈10pc,nH ≈10cm-3, fH2 ≈10-3, T>=200K
2. Rapid 3-body H2 formation for nH>1010cm-3
strong cooling supersonic inflow.
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The initial conditions for primordial star formation
Chemical Composition Trace H2 formation: H + e— H— +
H + H— H2 + e—
Tmin ~= 200 K, ncrit ~= 104 cm-3
MBE = 380Msun cs=1.2 km/s
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The initial conditions for primordial star formation
Tmin ~= 200 K, ncrit ~= 104 cm-3
MBE = 380Msun cs=1.2 km/s
Centrally concentrated cloud, inefficient cooling -> quasi-hydrostatic contractionDensity structure: ~self-similar, r -2.2
More chemistry: at high density >108cm-3
H+H+H -> H2 + Hefficient cooling -> dynamical collapse
Rotation: core forms from mergers andcollapse along filaments: expect J>0fKep vcirc / vKep ~ 0.5
Abel, Bryan, Norman (2002)
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The Accretion Rateand Formation Timescale
Density structure: self-similar, r -k, k≈2.2~singular polytropic sphere in virial and hydrostatic equilibrium P = K , =1.1
r -k
Ripamonti et al. 2002
Om
ukai & Nishi 1998
Abel et al. 2002
.Accretion rate: m*= * m / tff(m) f(m,K)(Tan & McKee 2004)
K=1.9x1012(T/300K)(nH/104cm-3)-0.1cgsK’=K/ 1.9x1012cgs*=1.4 (Hunter 1977)
m*(t=2Myr) ≈ 2000Msun
.“Isentropic Accretion”
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Collapse to a Disk
Geometry of Streamlines
Conserve J during free-fallinside sonic point (Ulrich 1976)
rd = f2Kepr0 3.4 (M/Msun)9/7 AU
Anticipate accretion drivenby large scale grav. instabilitiesand local gravitational viscosity (Gammie 2001)
viscosity = cs h, <0.3
fragmentationtcool< 3-1
STAR
DISK
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Toomre stability parameter
<1 -> unstable>1 -> stable
orbital angular velocity
soundspeed
surfacedensity
Are disks are stable with respect to fragmentation?
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Disk Models
Surface density
Thickness
Ionization
TemperatureTc , Teff
Toomre Q
=0.3
m*
.17x10-3Msun/yr 6.4x10-3Msun/yr 2.4x10-3Msun/yr
(Shakura &Sunyaev 1973; Tan & McKee 2004; Tan & Blackman 2004)
Disks arestable
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One zone model:follow energy ofprotostar as it
accretes
Evolution of the Protostar depends on Accretion Rate
Deuterium burning for Tc>106KStructural rearrangement after tKelvin
Eddington model for Solve for r*(m*), until reach main sequence
Assume polytropic stellar structure and continuous sequence of equilibria
(Stahler et al. 1980; Palla & Stahler 1991; Nakano 1995; Behrend & Maeder 2001; Omukai & Palla 2001, 2003; Tan & McKee 2003)
r*
m*
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Evolution of the Protostar
Initial conditionm* = 0.04 Msun
r* = 14 Rsun
(Ripamonti et al. 02)
Protostar is large (~100 Rsun) until it is older than tKelvin
Contraction to Main SequenceAccretion along Main Sequence
Comparison with Stahler et al. (1986), Omukai & Palla (2001)
Photosphere
Accretion Shock
Main Sequence(Schaerer 2002)
:Radius
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Evolution of the Protostar:Luminosity
Boundary Layer
Accretion Disk
Internal
Total
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Evolution of the Protostar :IonizingLuminosity
Spherical, fKep=0
fKep=0.5Total
Internal
Boundary Layer
Accretion Disk
Main Sequence(Schaerer 2002)
Spectrum depends on initial rotationc.f. Omukai & Palla (2003)
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Growth of the HII RegionBalance ionizing flux vs recombinations and infall
Find stellar mass at breakoutrHII = rg
polar; equatorial
Breakout mass vs rotation
Infall likely to be suppressedfor rHII>rg , where vesc=ci
HII Region
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1 in HI around HII region :
Photons diffuse in freq. and spaceNormalize J to appropriate Velocity field: Voigt profile D ;Line profile: damping wings
Ly- and FUV Radiation Pressure
L
Escape after n scatterings, or 2 photon decayfreq. shift ; mean free path at escapediffusion scale must equal size of region,
and
total path length of photons is n1/2L so mean intensity boosted by factor
(Neufeld 90) : Evaluate NH from harmonic mean of sightlines from star
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Mass Limits vs. Core Rotation
Breakout in polar direction
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Disk Photoevaporation
Weak wind case:
for zero age main sequence
.
Hollenbach et al. 94
Equate with mass accretion rate
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Mass Limits vs. Core Rotation
Disk Photo-evaporation
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When does accretion end?
Ly- Radiation Pressure
Ionization
Disk Evaporation
Hydromagnetic Outflows
Declining accretion versus increasing feedback
m* >~ 20-30 Msun (polar)
m* >~ 100 Msun
m* ~ 100-200 Msun
m* >~ 100-500 Msun
Tan & Blackman (2004)
McK
ee &
Tan
, in
prep
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ConclusionsConvergent initial conditions for star formation: set by H2 cooling
Accretion rate + semi-analytic model for protostellar evolution:large (~AU) protostar, contracts to main sequence for m*>30M
This is when feedback processes start to become important
Feedback processes depend on core rotation and are complicated:Gradual reduction in SF efficiency because of ionizing and radiationpressure feedback for m*>30M. Final mass, ~100-200M, likely to be set by ionizing feedback on the accretion disk
Implications of massive star formation in each mini-halo?Are low-mass zero metallicity stars possible?How effective is external feedback?Is this mode of star formation inevitable in all zero metal DM halos?Are these the seeds of supermassive black holes?
Mostly thermal pressure support + slow cooling -> no fragmentation -> single ~massive star in each mini-halo
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Evolution of accretion and outflow rates
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Vertical structure of disk
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Evolution of ionizing luminosity with varying mdot
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Evolution of Lyman-Werner photon luminosity
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overview
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Radial profile of disk
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Radial profile allowing for varying mdot
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Comparison: then and now
The first stars The latest starsH2 cooling, T~200K CO/dust cooling, T~10-20Kns ~104cm-3; M~300Msun ns ~106cm-3; dn/dM ~ M-2
thermal pressure nonthermal (B) pressure, turbulent-> single, isolated stars -> fragmentation to star cluster
ionization, Ly-, Rad.pressure on dustMHD outflows? MHD outflows
Combined feedback of many stars
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Initial Conditions for Star Formation from Abel, Bryan, Norman 02
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core rotation (ABN)
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