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Seminario Italia-Giappone
of the First Stars
Kazuyuki Omukai
(NAO Japan)
First Stars:
proposed as an origin of heavy elements
Sun 2%, metal poor stars 0.001-0.00001%
Cause of early reionization of IGM
te=0.17 zreion=17 (WMAP)
Depend on mass /formation rate of first stars
Let’s study their formation process !
Before the First Stars
Cosmological initial condition (well-defined)
Pristine H, He gas, no dusts, no radiation field
(except CMB), CR
simple chemistry and thermal process
No magnetic field (simple dynamics)
After the First Stars
Feedback (SN, stellar wind) turbulent ISM
metal /dust enriched gas
radiation field (except CMB), CR
complicated microphysics
magnetic field
Formation of First Objects:
condition for star formation
 Hierarchical
small objects form
 Condition for star
radiative cooling is
necessary for further
contraction and star
Tegmark et al. 1997
First Objects (3s)
z~30, M~106Msun
cool by H2
Microphysics of Primordial Gas
Radiative cooling rate
In primordial gas
Atomic cooling only
effective for T>104K
Below 104K, H2 cooling
is important
H2 formation
(H- channel: e
H + e -> H- + g
H- + H -> H2 + e
Efficient cooling for T>1000K
Simulating the formation of first objects
Yoshida, Abel, Hernquist & Sugiyama (2003)
ab initio calculation is already possible !
Road to the First Star Formation 1
1. Formation
of the First Object
Road to the First Star Formation 2
2. Fragmentation
of the First Objects
Fragmentation of First Objects
3D numerical simulation is getting possible
3D similation
(Abel et al. 2002,Bromm et al. 2001)
filamentary clouds
(Nakamura & Umemura 2001)
Typical mass scale of
Dense cores
Bromm et al.. 2001
a few x 102-103Msun
no further fragmention
These cores will collapse and form protostars eventually.
Road to the First Star Formation 3
3. Collapse of Dense Cores:
Formation of Protostar
60% known
Pop III Dense Cores to Protostars: Thermal Evolution
cooling agents:
H2 lines
(log n<14)
H2 continuum
becomes opaque
at log n=16
H2 dissociation
(K.O. & Nishi 1998)
Temperature evolution
approximately, g =d log p/d log n= 1.1
Pop III Dense Cores to Protostars: Dynamical Evolution
self-similar collapse
up to
(K.O. & Nishi 1998)
protostar formation
state 6; n~1022cm-3,
Mstar~10-3Msun Tiny Protostar
(~Pop I protostar )
3D simulation for prestellar collapse
Abel, Bryan & Norman 2002
The 3D calculation
has reached
(radiative transfer
needed for higher
density; cf. n~1022cm-3
for protostars)
Overall evolution is
similar to the 1D
The collapse velocity
is slower.
(why? the effect of
rotation, initial
condition, turbulence)
Road to the First Star Formation 4
4. Accretion of ambient gas and
Relaxation to Main Sequence Star
25% known
Density Distribution at protostar formation
Density around the primordial protostar is higher
Than that around prensent-day counterpart.
(For hot clouds, the density must be higher to overcome the
stronger pressure and form stars.)
This difference affects the evolution after the protostar formaition
via accretion rate.
Mass Accretion Rate
After formation, the protostars grow in
mass by accretion.
The accretion rate is related to density distribution
(the temperature in prestellar clumps):
3/ 2
M  M Jeans / t ff  cs / G  T
Pop III T~300K
Mdot ~ 10-3 – 10-2Msun/yr
Pop I
Mdot ~ 10-6 - 10-5Msun/yr
The accretion rate is very high
for Pop III protostars
Protostellar Evolution in Accretion Phase
Protostellar Radius
M = 8.8, 4.4, 2.2, 1.110 3 M  / yr
Nuclear burning is delayed by accretion.
(K.O. & Palla 2003)
(H burning via CN cycle at several x10Msun)
Accretion continues in low Mdot cases, while the stellar
wind prohibit further accretion in high Mdot cases.
Critical accretion rate
Total Luminosity (if ZAMS)
Ltot  LZAMS  GM  M / RZAMS
Exceeds Eddington limit
if the accretion rate is larger than
M crit = 4c(1  LZAMS / LEdd ) RZAMS /  es
 4 10 M  / yr
In the case that Mdot > Mdot_crit, the stars cannot reach the ZAMS
structure with continuing accretion.
How much is the Actual Accretion Rate ?
From the density distribution
around the protostar…
Abel, Bryan, & Norman
Protostellar Evolution for ABN Accretion Rate
Evolution of radius
under the ABN accretion rate
The protostar reaches
ZAMS after Mdot
decreases < Mdot_crit.
Accretion continues….
The final stellar mass
will be 600Msun.
Pop I vs Pop III Star Formation
Pop I core
Mstar : 10-3Msun
Mclump: >0.1Msun
Mdot: 10-5Msun
With dust grains
Pop III core
Mstar : 10-3Msun
Mclump : >103Msun
Mdot : 10-3Msun
No dust grain
Massive stars (>10Msun)
are difficult to form.
Accretion continues.
Very massive star formation
a 2nd generation star found !
Most iron-deficient star
HE0107-5240 [Fe/H]=-5.3
Iron less than
10-5 of solar;
Second Generation
Low-mass star
What mechanism
causes the transition
to low-mass star
formation mode?
Christlieb et al. 2002
Key Ingredients in 2nd Generation
Star Formation
Metal Enrichment
UV Radiation Field from pre-existing stars
Density Fluctuation created by SN blast
wave, stellar wind, HII regions
Metals from the First SNe
Heger, Baraffe, Woosley 2001
Type II SN
Pair-instability SN
Metals and Fragmentation scales
K.O.(2000), Schneider, Ferrara, Natarajan, & K.O. (2002)
Formation of massive fragments continues
until Z~10-4Zsun (If radiation not important)
For higher metallicity, sub-solar mass
fragmentation is possible.
Radiation pressure onto dusts
if d>es, radiation
pressure onto dust shell
is more important.
massive SF
 This occurs ~0.01Zsun
 For Z<0.01Zsun
Accretion is not halted
Metals and Mass of Stars
Massive frag.
Low-mass frag. possible
Accretion not halted
Massive stars
& massive
Accretion halted by
dust rad force
Effects of UV Radiation Field
Star Formation in Small Objects (Tvir < 104K)
(K.O. & Nishi 1999)
Only one or a few massive stars can
photodissociate entire parental objects.
Without H2 cooling, following star formation is
Only One star is formed at a time.
FUV radiation effect on fragmentation scale
Star formation in large objects (Tvir>104K)
Evolution of T in the prestellar collapse
Fragmentaion scale vs 5UV intensity
radiation: Jn=W Bn(10 K) from massive PopIII stars
K.O. & Yoshii 2003
H2 cooling clumps
critical value
(logW < -15)
W<Wcrit H2 formation,
cooling sun
crit no H2
(logW > -15)
(Lyα –– H- f-b
In starburst of large objects, subsolar mass Pop III
scale decreases for stronger radiation
can be formed.
Effects of SN blast wave
(Wada & Venkatesan 2002; Salvaterra et al. 2003)
SNe of metal-free stars
(Umeda & Nomoto 2002)
SN II (10Msun-30Msun; 1051 erg)
pair instability SN
(150Msun-250Msun; 1053erg)
Shell formation by blast
fragmentation of the shell
low-mass star formation?
Bromm, Yoshida, & Hernquist 2003
Typical mass scale of the first stars
is very massive ~102-3Msun,
because of
large fragmentation,
(ii) continuing accretion at large rate
However, the conclusion is still rather qualitative.
Formation of the second generation
of stars is still quite uncertain.
Metallicity/ radiation can induce the transition from
massive to low-mass star formation mode.