Kazuyuki Tanaka Graduate School of Information Sciences Tohoku University, Sendai 980-8579, Japan
Kazuyuki Omukai (NAO Japan) Formation of the First Stars Seminario Italia-Giappone.
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Transcript of Kazuyuki Omukai (NAO Japan) Formation of the First Stars Seminario Italia-Giappone.
Kazuyuki Omukai (NAO Japan)
Formation of the First Stars
Seminario Italia-Giappone
First Stars:
proposed as an origin of heavy elements
Sun 2%, metal poor stars 0.001-0.00001%
Cause of early reionization of IGM
ezreion=17 (WMAP)
Let’s study their formation process !
Depend on mass /formation rate of first stars
Before the First StarsCosmological 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 ISMmetal /dust enriched gas radiation field (except CMB), CR
complicated microphysics magnetic field MHD
SIMPLE
COMPLICATED
Hierarchical clustering
small objects form earlier
Condition for star formation
radiative cooling is necessary for further contraction and star formation
First Objects (3) z~ 30, M ~ 106Msun
Tvir ~ 3000K cool by H2
Formation of First Objects: condition for star formation
Tegmark et al. 1997
Atomic cooling only effective for T>104K
Below 104K, H2 cooling is important
H2 formation
(H- channel: e catalyst)
Radiative cooling rate In primordial gas
H + e -> H- + H- + H -> H2 + e
Efficient cooling for T>1000K
Easy Microphysics of Primordial Gas
Simulating the formation of first objects
ab initio calculation is already possible !
Yoshida, Abel, Hernquist & Sugiyama (2003)
600h-1kpc
1. Formationof the First Object
Road to the First Star Formation 1
95%known
2. Fragmentationof the First Objects
Road to the First Star Formation 2
50%known
3D similation (Abel et al. 2002,Bromm et al. 2001)filamentary clouds (Nakamura & Umemura 2001)
Bromm et al.. 2001
Typical mass scale of fragmentation;
Dense cores
a few x 102-103Msun
no further fragmention
Fragmentation of First Objects
3D numerical simulation is getting possible
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
( K.O. & Nishi 1998)
Pop III Dense Cores to Protostars: Thermal Evolution
cooling agents: H2 lines (log n < 14) H2 continuum (14-16) becomes opaque
at log n=16 H2 dissociation (16-20)
Temperature evolution
approximately, =d log p/d log n= 1.1
( K.O. & Nishi 1998)
self-similar collapse
up to n~1020cm-3
protostar formation
state 6; n~1022cm-3, Mstar~10-3Msun
( ~Pop I protostar )
Pop III Dense Cores to Protostars: Dynamical Evolution
Tiny Protostar
3D simulation for prestellar collapse
The 3D calculation has reached n~1012cm-3
(radiative transfer needed for higher density; cf. n~1022cm-3 for protostars)
Overall evolution is similar to the 1D calculation.
The collapse velocity is slower.
(why? the effect of rotation, initial condition, turbulence)
Abel, Bryan & Norman 2002
Road to the First Star Formation 4
4. Accretion of ambient gas andRelaxation to Main Sequence Star
25% known
Density Distribution at protostar formation
(For hot clouds, the density must be higher to overcome the stronger pressure and form stars.)
Density around the primordial protostar is higherThan that around prensent-day counterpart.
This difference affects the evolution after the protostar formaitionvia accretion rate.
Mass Accretion Rate
After formation, the protostars grow in mass by accretion.
2/33 // TGctMM sffJeans
The accretion rate is related to density distribution(the temperature in prestellar clumps):
Pop III T~300K Mdot ~ 10-3 – 10-2Msun/yr
Pop I T~10K Mdot ~ 10-6 - 10-5Msun/yr
The accretion rate is very high for Pop III protostars
Protostellar Evolution in Accretion Phase Protostellar Radius yrMM /101.1 ,2.2 ,4.4 ,8.8 3
3 a, ZAMS
3b、 expansion
2, KH
contr.
1、 adiabatic phase
tKH >tacc
Nuclear burning is delayed by accretion.
(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.
(K.O. & Palla 2003)
Critical accretion rate
Total Luminosity (if ZAMS)
ZAMSZAMStot RMGMLL /
Exceeds Eddington limit if the accretion rate is larger than
yrM
RLLcM esZAMSEddZAMScrit
/104
/)/1(43-
In the case that Mdot > Mdot_crit, the stars cannot reach the ZAMS structure with continuing accretion.
Abel, Bryan, & Norman (2002)
How much is the Actual Accretion Rate ?
From the density distribution around the protostar…
Protostellar Evolution for ABN Accretion Rate
The protostar reaches ZAMS after Mdot decreases < Mdot_crit.Accretion continues….The final stellar mass will be 600Msun.
Evolution of radius under the ABN accretion rate
Pop I vs Pop III Star Formation
Pop I coreMstar : 10-3Msun
Mclump: >0.1Msun
Mdot: 10-5Msun
With dust grains
Pop III coreMstar : 10-3Msun
Mclump : >103Msun
Mdot : 10-3Msun
No dust grainMassive stars (>10Msun)are difficult to form.
Accretion continues.Very massive star formation (100-1000Msun)
a 2nd generation star found !
Iron less than
10-5 of solar;
Second Generation
Low-mass star ~0.8Msun
What mechanism causes the transition to low-mass star formation mode?
Most iron-deficient star HE0107-5240 [Fe/H]=-5.3
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
Type II SN 8-25Msun
Pair-instability SN 150-250Msun
Heger, Baraffe, Woosley 2001
PISNSN II
Metals and Fragmentation scales
Formation of massive fragments continues until Z~10-4Zsun (If radiation not important)
For higher metallicity, sub-solar mass fragmentation is possible.
K.O.(2000), Schneider, Ferrara, Natarajan, & K.O. (2002)
Radiation pressure onto
dusts
ifd>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
0 Zsun10-5Zsun 10-2Zsun
Massive frag. Low-mass frag. possible
Accretion not haltedAccretion halted by dust rad force
Massive stars Low-mass & massivestars
Low-massstars
Effects of UV Radiation Field
Only one or a few massive stars can photodissociate entire parental objects.
Without H2 cooling, following star formation is inhibited.
(K.O. & Nishi 1999)
Star Formation in Small Objects (Tvir < 104K)
Photodissociation
Only One star is formed at a time.
FUV radiation effect on fragmentation scale
log(W)=-15 ; critical value
W < Wcrit H2 formation, and cooling
W>Wcrit no H2
( Lyα –– H- f-b cooling)
Evolution of T in the prestellar collapse
radiation : J=W B(105K) from massive PopIII stars
Star formation in large objects (Tvir>104K)
Fragmentation scale H2 cooling clumps ( logW < -15 )
Mfrag~2000-40Msun
Atomic cooling clumps (logW > -15)
Mfrag~0.3Msun
Fragmentaion scale decreases for stronger radiation
Fragmentaion scale vs UV intensity
In starburst of large objects, subsolar mass Pop III Stars can be formed.
K.O. & Yoshii 2003
Effects of SN blast wave
SNe of metal-free stars (Umeda & Nomoto 2002) SN II (10Msun-30Msun; 1051 erg) pair instability SN
(150Msun-250Msun; 1053erg)
Shell formation by blast wave
fragmentation of the shell
low-mass star formation?
(Wada & Venkatesan 2002; Salvaterra et al. 2003)
Bromm, Yoshida, & Hernquist 2003
ConclusionTypical mass scale of the first stars is very massive ~102-3Msun,
because of
(i) 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.