Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Evidence for GR Induced Explosions inSupermassive Population III Stars
Chris Nagele; Hideyuki Umeda, Koh Takahashi, Takashi
Yoshida, Kohsuke Sumiyoshi
University of Tokyo

November 11, 2019

Summary

Background

GR Instability

Outline1

Background2

GR Instability

3

Method: HOSHI

4

Method: Radhyd

5

Results

6

Summary

Method: HOSHI

Method: Radhyd

Results

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Summary

The Problem
We observe high mass, high z quasars
How were supermassive black holes (SMBHs) formed at such
early times?
Image credit: Dr. Sunmyon CHON (Tohoku University)

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Summary

Direct Collapse Scenario
A large cloud of gas collapses into a supermassive star (SMS)The star continues to accrete during its lifetime
The star collapses into an intermediate mass black hole(IMBH)

Gas cloud−→M ∼ 107−8 M(Bromm Loeb 2002)

SMS−→
IMBH
M ∼ 104−5 M
M ∼ 105 M(EHT Collab. 2019)

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Direct Collapse Scenario
M ∼ 107 M , T ∼ 104 K metal free gas halo
Nearby source of Lyman Werner photons suppresses H2formation
Resulting high virial temperature prevents fragmentation
Rare but no impossible

No Stars Yet

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Summary

IMBH Growth

Once an intermediate mass black hole is formed, can it growinto a supermassive black hole (SMBH)?

Large scale studies suggest that IMBHs can grow at theEddington rate (e.g. Di Matteo+ 2012)
Studies incorporating smaller scale physics find far less growth(e.g. Latif+ 2018)As of yet no consensus

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

IMBH Growth
Di Matteo+ 2012

Latif+ 2018

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Chen+ 2014
General Relativistic Instability Supernova of aSupermassive Population III Star
Evolved Pop III main sequence progenitor until the point ofcollapse
Simulated the collapse in 1D and 2D

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Chen+ 2014

They found that M = 55000M , (M(He) = 31000M )collapses due to GR instability
Early collapse leads to rapid 4 He burning and an explosion
with energy E = 6 − 9 ∗ 1054 ergs
Narrow mass range for explosion’Ideal Case’ in the context of the direct collapse scenario

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Instability Condition: Two Approaches (Butler+ 2017)
A completely radiative gas (Γ = 4/3) is unstable
Gas pressure stabilizes the star
Approach 1

Approach 2

Do not assume polytropic EOS

Assume polytropic EOS

Assume constant entropy withPgas as a perturbation to PradIf Γ < 4/3, the star will beunstable

Causes star to act like n = 3
polytrope

Stability condition:

Stability condition:

ρc < ρcritρcrit = 1.996∗1018 (M 7/2) g cm−3Mβ1.12 ∗ GM>
6r2

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Pop III Stellar Evolution

HOSHI: Takahashi+ 2016, 2018
Population III 1D stellar evolution
Large reaction network (212 isotopes)
Includes rotation (deformation, rotational instabilities)

Summary

Background

GR Instability

ρ-T Diagrams

Method: HOSHI

Method: Radhyd

Results

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

TOV Correction

Tolman–Oppenheimer–Volkoff equation
General relatvistic version of hydrostatic equation
Becomes important for stars with M > 1000M
geffective =−GmP
4πr 3 P
2Gm −1(1+)(1+)(1 +)222r
ρ0 cmcrc 2

geffective ∼= g (1 +P
4πr 3 P
2Gm++)
22
ρ0 cmcrc 2

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Summary

Collapse
Radhyd: Yamada 1997, Takahashi+ 2018
1D-spherical general-relativistic Lagrangian hydrodynamic code
Updated in Takahashi+ 2018 to include nuclear networks
49 and 153 isotope networks
Neutrino cooling
Pair Instability SN (100 MHe)

Collapse to Black Hole (200 M He)

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Drawbacks

HOSHI does not conserve energy during collapse
HOSHI does not have full GR treatment
Radhyd may overestimate He burning, especially for slow
collapse
Radhyd is not always stable on timescales τKH
When should we switch between codes?

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Explosion (Preliminary)55, 000M∆M(He) = 2762M
Eexplosion = 1.05 ∗ 1055 ergs
Long ’explosion’ time
Starts with adiabatic collapse of entire star

Results

Summary

Background

GR Instability

Method: HOSHI

Explosion (Preliminary)55, 000M∆M(He) = 2762M
Eexplosion = 1.05 ∗ 1055 ergs
Long ’explosion’ time

Method: Radhyd

Results

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Mass Range (Preliminary)
Chen+ 2014 only found 1 exploding model
Many GR unstable models
Many exploding models
Not necessarily overlapping

Results

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Parameters (Preliminary)He abundance during code transfer
Mass
GR Stability

Results

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

Summary + Future Work

Summary
Chen+ 2014 found a GR induced supernova for a 55000 M
Pop III star.
Preliminary results show some evidence that the explosion
existsFuture WorkWhen should we switch between codes?
Implementing advection heating
Establishing lower mass limit
Reproducing Fuller+ 1986

Summary

Background

GR Instability

Method: HOSHI

Method: Radhyd

Results

SummaryWorks Cited

Butler, S. P. Lima, A. R. Baumgarte1, T. W. Shapiro, S. L. 2017, MNRAS 477, 3694
Chandrasekhar, S. 1964, PhRvL, 12, 114
Chen, K.-J., Heger, A., Woosley, S., et al. 2014, ApJ, 790, 162
Di Matteo, T., Khandai, N., DeGraf, C., et al. 2012, ApJL, 745, L29
Fuller, G. M., Woosley, S. E., Weaver, T. A. 1986, ApJ, 307, 675
Latif, M. A., Volonteri M., Wise J. H. 2018, MNRAS, 476, 4Mortlock, D. J., Warren, S. J., Venemans, B. P., et al. 2011, Natur, 474, 616
Takahashi, K., Yoshida, T., Umeda, H., Sumiyoshi, K., Yamada, S. 2016, MNRAS,456, 1320
Takahashi, K. Umeda, H. Yoshida, T. et al. 2018, APJ 857:111
Wu, X.-B., Wang, F., Fan, X., et al. 2015, Natur, 518, 512