Contents Star formation in the early universe - from big bang - - PowerPoint PPT Presentation
Ancient Universe with GRB, Kyoto, 2010 T HE F IRST S TARS Naoki Yoshida Contents Star formation in the early universe - from big bang ripples to protostars The fate of very massive stars Low-metallicity star formation References: NY,
Naoki Yoshida Ancient Universe with GRB, Kyoto, 2010
✦ Star formation in the early universe
✦ The fate of very massive stars ✦ Low-metallicity star formation
References: NY, Omukai, Hernquist, 2008, Science Bromm, NY, McKee, Hernquist, 2009, Nature (review) Ohkubo, Umeda, Nomoto, NY, Tsuruta, 2009, ApJ Umeda, NY, Nomoto, Ohkubo, Sasaki, Tsuruta, 2009, JCAP Omukai, Hosokawa, NY in prep.
From the ancient universe
A massive star’s death just 600 million years after the Big Bang! GRB 090423 @z=8.3
In the backyard...
Frebel et al. 2005
The Standard Cosmology
CMB + LSS + SNe tell us about the initial state of the universe, and the energy content now and then precisely.
An ab initio approach is possible
density fluctuations
length scale
Primordial Star Formation
A ‘Simple’ Problem
The Physics - understood Gravity, hydrodynamics, atomic/molecular processes 14 species (H, He, D, ions) ~50 important reactions + many radiative processes Density evolution to ~1021/cc The Initial Condition - “observed” Cosmologically determined; CDM model dark matter + hydrogen-helium gas + CMB In the Dark Age...
In the beginning, there was a sea of light elements and dark matter ....and some ripples.
From diffuse gas to protostar
adiabatic contraction H2 formation line cooling
(NLTE)
loitering
(~LTE)
3-body reaction Heat release
molecular line collision induced emission
T [K]
104 103 102
number density
continuum and dissociation A proto-star (hydrostatic core) The Physics
The Omukai diagram
MJ~1000Msun
A complete picture
process of a primordial protostar. Dynamic range 1013 Resolving planetary scale structures in a cosmological volume!
NY, Omukai, Hernquist Science 2008
An early universe “experiment”
Composition Velocity The central 0.01 Msun core is a seed for the subsequent formation of a massive star.
atomic core fully molecular
A hyper-accreting protostar
hydrostatic core
“hot” infalling gas
T ~ 500-10000K The gas mass accretion rate is estimated to be dM/dt = 0.01-0.1 Msun/yr. This is enough to make a 10-100 Msun star within 1000 years!
NY, Omukai, Hernquist, Abel, 2006, ApJ
Protostellar evolution
MZAMS ~ 100 Msun dM/dt = 0.01-0.1 Msun/yr (time dependent)
H burning starts
Primordial Star Formation
the onset of collapse.
surrounding the protostar = Very large accretion rate
= accretion continues
Bromm+99 Nakamura-Umemura 01 Omukai-Palla 03 NY+ 06, 08 Omukai-Palla 03 Tan-McKee 04 Hosokawa’s poster
A stick in our throat
Very massive Pop III... Models and numerical simulations suggest that the first stars were rather massive, > 100 Msun However, there is no indication, no single evidence that pair-instability supernovae contributed Galactic chemical evolution.
From N. Tominaga
Observed metal-poor stars
expelling metals (Madau&Rees01, Ohkubo+09)
star surveys? (Karlsson+08)
star formation (Johnson&Bromm06, NY, Omukai 07)
expelling metals (Madau&Rees01, Ohkubo+09)
star surveys? (Karlsson+08)
star formation (Johnson&Bromm06, NY+07)
Model accretion rate
0.1 0.01 0.001 1 10 100 1000 Msun
Msun/year
Accretion rates from
Y06, 1st
Y07, 2nd gen.
Ohkubo, Umeda, Nomoto, NY, Tsuruta, ApJ, 2009
Hyper Massive Star
1 10 Mass 100 1000 Stellar radius
Hydrogen in the core exhausted
Ohkubo, Umeda, Nomoto, NY, Tsuruta 2009
P a i r I n s t a b i l i t y S N Direct collapse
NY, Oh, Kitayama, Hernquist (2007) ApJ
Star formation in a reionized gas Temperature profile
HD cools the gas! CMB plays a role!
PopIII.2 after reionization
PopIII.2 after reionization
1st star
with HD cooling
NY, Omukai, Hernquist (2007, ApJL, 667, 117)
Primordial stars formed by HD cooling are not very massive because of low-T ~ low dM/dt Hydrogen-burning starts at M ~ 30Msun MZAMS ~ 40Msun Mcloud ~ 40Msun
mass Protostellar radius
NY, Omukai, Hernquist 2007, ApJL Nakamura & Umemura 2003 Nagakura & Omukai 2005 McGreer & Bryan 2008
Stars formed by HD cooling “conventional”
Is there a “critical metallicity” for cloud fragmentation ? If so, what determines it ?
Bromm et al. cooling by C, O @low-density Omukai, Schneider cooling by dust @high density
vs.
Toward a direct simulation
Chemistry and radiative transfer in a gas with heavy elements and dust : １ Cooling by CI, CII, OI ２ Dust thermal emission ３ Molecular cooling by H2O, OH, CO ４ New cooling rates for H2, HD
Tdust determined by the thermal balance: 4 T 4 = Lgr (gas -> dust) Tdust temperature evolution Next talk by Kaz!
密度 温 度 H2 HD H2O Tdust
Oxygen chemistry : Z=-5
密度 [/cc] H2O OI OH
O + H -> OH + H2 + OH -> H2O + H
For Z=-5, Rapid cooling by dust at high density (n~1014) leads to core fragmentation. Fragment mass ~ 0.1 Msun 5AU
massive
supermassive blackholes at high-z
cloud fragments by dust cooling,
Turk+ (2009) simulation: the core breaks up at n ~ 1012 /cc. Formation of massive (not very massive) star pair with ~ tens Msun. 1case over 5 samples. A large fraction produces PopIII binaries ?
Case with a fitting function (a la Turk et al.)
A direct comparison of density evolution With 3D radiative transfer a factor of ~100 at a given time
Radiative transfer effects
Molecular line escape probability Velocity gradient
Similarly important for continuum
The escape probability is not just a function of density; it varies with temperature (via Doppler width) and with velocity gradients, hence is highly time- and direction-dependent.
NY, Omukai in prep