Formation of Primordial Protostars The Final Chapter

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Formation of Primordial Protostars The Final
Chapter
First Stars III Santa Fe, July 19, 2007
Naoki Yoshida Department of Physics Nagoya
University
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Turbulence is important

• The first flight Nagoya-Tokyo was
• cancelled due to a big typhoon. (Turbulence)
• The second flight from Tokyo passed
• over the wakes of the typhoon and it terribly
• shaked us. (Real free-fall.)
• The last flight from LA to Albuquerque
• returned to LA because of some trouble
• in anti-icing system.
• (chemistry, phase-transition).

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Contents

• Thermal Evolution of a Primordial Gas
• - Physics at high densities (cooling,
  chemistry, radiation)
• - Chemo-thermal instability
• - Accretion shock and proto-stellar
  evolution
• NY, Omukai, Hernquist, Abel (2006, ApJ,
  652, 6)
• NY et al. in preparation
• Formation of Primordial Stars in a Reionized Gas
• - Early HII/HeIII regions, H2/HD cooling
• - 10-40Msun primordial stars and
  hypernovae/GRB
• Yoshida (2006, NewA, 50, 19)
• NY, Oh, Kitayama, Hernquist (2007, ApJ,
  663, 687)
• NY, Omukai, Hernquist (2007,
  arxiv/0706.3597)
• Primordial Stars and Dark Matter

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An ab initio simulation of theformation of
primordial stars
The Initial Condition
?CDM model, Gaussian random density field dark
matter hydrogen-helium gas CMB
Gravity, hydrodynamics, atomic/molecular
processes
14 species, equilibrium, non-eq. e, H, H, H-,
H2, H2, He, He, He, D, D, D-, HD, HD 50
reactions many radiative processes
Density evolution to 1021 cm-3
Initial density field
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Self-gravitating cloud
0.3Mpc
5pc
0.01pc
A new born proto-star with T 20,000K
r 10 Rsun!
Fully-molecular core
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How did we achieve thisextreme resolution ?
Physics, physics, and more physics, (and solving
troubles.)
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Thermal evolution of a primordial gas
adiabatic phase
collision induced emission
104
H2 formation line cooling (NLTE)
T K
3-body reaction heat release
opaque to cont. full-scale dissociation
103
loitering (LTE)
opaque to molecular line
adiabatic contraction
102
number density
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We have finally done it!
proto-star
VELOCITY PROFILE
104
Real gas effects, Pressure ionization
Effective Equation of State from a
first-principle calculation
T K
103
102
number density
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Raditive Transfer 1molecular lines
1010-14

• Example) J6?4 transition at T1000 K

For ? gt0.1, the cloud core becomes optically
thick to H2 lines?and then line cooling is
inefficient (t4,6 1 for Lcore0.0001pc)
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The Sobolev length and escape probability
V thermal dV/dr
Lsobolev , ? ?
Lsobolev We compute the Sobolev lengths from
local velocity gradients dVx/dx, dVy/dy, dVz/dz
Physically well-motivated. It can be applied to
more general problems than an assumed EoS can
be.
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Net molecular cooling rate
?thick ? ?escape nk,l A k,l h?
3D calculation
optically thick / thin
CIE cooling
Ripamonti???? ????(???) ????
Omukai98 1D full RT
8 10 12 14
16 log (n)
NY, Omukai, Hernquist, Abel (2006)
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Cooling by continuumCollision Induced Emission
1014-17
h??
continuum emission
During collisions, a collision pair acts as a
super-molecule, generating an induced electric
dipole.
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Optically thin CIE cooling rate
2h?3 c2
??(?) ? nH2 exp(-h?/kT)
H2-H2 (Borysow et al. 2001) H2-He (Jorgensen et
al 2000) _at_ n 1013 - 1017
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Radiative Transfer 2continuumGrey transfer and
the Planck opacity
1015-18
A standard table from Lenzuni et al. (1991) An
updated version from Mayer Duschl (2005)
Lenzuni Duschl
? For n 1016 /cc
T K
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Equilibrium chemistry the Saha equations
gt1016
A set of coupled equations including the energy
equation
are solved self-consistently
H-H
H-H2
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Cooling/heating processes
CIE
Chemical reaction
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Cosmological Simulations
Standard ?CDM model Multi-level zoom-in
technique final mass resolution 30Mmoon final
spatial resolution Rsun
Hydro, eq/neq-chemistry , radiative processes,
etc. etc.
NY, Omukai, Hernquist, Abel (2006)
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Chemo-thermal instabilityNumerical results
2 1.5 1 0.5 0
growth parameter
tff/tg becomes larger than 1, but always below
2. Perturbation growth and gravitational
collpase occur on a similar time scale. Although
the thermal instability occurs, the cloud does
not fragment to multiple objects - Collapse is
just accelerated.
tfree fall / tgrowth

• 9 10 11 12

log (n)
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Primordial proto-stara tiny seed in a large
cloud
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Accretion rate and proto-stellar evolution
NY, Omukai, Hernquist, Abel (2006)
dM/dt 0.1-0.001 Msun/yr MZAMS
60-100 Msun
We used the obtained accretion rate as an
input to the proto-stellar evolution calculation
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Primordial Star Formation in a ?CDM universe

• No fragmentation is observed during the
• prestellar collapse. The parent cloud is a
• single 300 Msun cloud at the center of
• a cosmological mini-halo.
• The collapsing cloud is stable against
• gravitational deformation, too.
• A tiny proto-stellar seed with mass
• 0.01 Msun is formed first.
• Proto-stellar evolution calculations give
• MZAMS 100 Msun

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Radiative feedback from the first starand 2nd
generation star formation
Effect of HD cooling
T K
TCMB at z16
NY, Oh, Kitayama, Hernquist (2007,
ApJ) Radiation-hydrodynamics calculation
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2nd generation primordial star
NY, Omukai, Hernquist (2007, astro-ph)
Primordial stars in a reionized gas are not very
massive
1st star
Hydrogen-burning starts at M30Msun MZAMS
40Msun Mcloud 40Msun
2nd. gen. star with HD cooling
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Implications

• 1 Parent gas cloud mass, accretion rate both
• substantially smaller than ordinary PopIII
  cases
• gt M lt 40 Msun (smaller mass at lower
  redshift)
• 2 The progenitor mass of the OBSERVED HMP
• stars is suggested to be 20-40 Msun
• (Iwamoto et al. 2005)
• 3 Also in a hypernova-GRB progenitor mass range
• Primordial stars formed during/after
  reionization
• likely trigger GRBs.

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First stars and dark matter
WDM simulation Gao Theuns, to appear in Science
Neutralinos as dark matter Spolyer et al. (2007)
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Primordial star formationThings to explore
further are

• Further evolution up to
• the adiabatic phase
• Accretion process,
• disk accretion
• 2. Feedback from the proto-star