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Pair-instability supernovae
From Woosley et al. (2002, 2007)
Woosley Lecture 19
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Mass Loss in Very Massive Primordial Stars
• Negligible line-driven winds (mass loss ~ metallicity1/2) (Kudritzki 2002)
• No opacity-driven pulsations (no metals)• Continuum-driven winds likely small contribution• Epsilon mechanism inefficient in metal-free stars
below ~1000 M (Baraffe, Heger & Woosley 2000)from pulsational analysis we estimate upper limits:– 120 solar masses: < 0.2 %– 300 solar masses: < 3.0 %– 500 solar masses: < 5.0 %– 1000 solar masses: < 12.0 %during central hydrogen burning
• Red Super Giant pulsations could lead to significant mass loss during helium burning for stars above ~500 M
8 – 11 M¯: uncertain situation
?
• M < M1 ' 8 M¯: No C ignition
• M > M2 ' 12 M¯: Full nondegenerate burning
• In between: ????
• Degenerate off-centre ignition
• Possibly O-Ne-(Mg?) white dwarfs (after some additional mass loss)
• With sufficient O-Ne core mass: continued burning and core collapse
Pair-instability supernovae
• He burning
• collapse and energy release
• + ! e+ + e-: 1 < 4/3
• Dynamical collapse, bounce, explosive burning (for M < 260 M¯)
• Dynamical collapse directly to black hole (for M > 260 M¯)
Pop. III stars, no mass loss
Possibly observed: SN 2006gy
Smith et al. (2007; ApJ 666, 1116)
Can very massive stars retain their masseven today?
The Pistol Star• Galactic star• Extremely high mass loss rate• Initial mass: 150 (?)• Will die as much less massive object
Pair instability
Helium core mostly convective and radiation a large part of the total pressure.~ 4/3. Contracts and heats up after helium burning. Ignites carbon burning radiatively
Above 1 x 109 K, pair neutrinos accelerate evolution. Contraction continues. Pair concentration increases. Energy goes into rest mass of pairs rather than increasing pressure, < 4/3. Contraction accelerates.
Oxygen and (off-center) carbon burn explosively liberating a large amount of energy. At higher mass silicon burns to 56Ni
The star completely, or partially explodes
Barkat, Rakavy and Sack (1967)
(M> 40 solar masses)
Helium stars, MHe = 2.2 – 8
degeneracy parameter
Nomoto and Hasimoto (1986; Prog. Part. Nucl. Phys. 17, 267)
Pair-Instability Supernovae
Many studies in literature since more than 3 decades, e.g.,Rakavey, Shaviv, & Zinamon (1967)Bond, Anett, & Carr (1984)Glatzel, Fricke, & El Eid (1985)Woosley (1986)
Some recent calculations:Umeda & Nomoto 2001Heger & Woosley 2002
63 M / Me 40
130 > M / Me 95
133 M / Me 63
260 M / Me 130
M 133 Me
M > 260 Me
Pulsational PairSupernovae
Pair instabilitySupernovae
Black holes
Rotation reduces thesemass limits!
Mass loss alters them.¯
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¯
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Light curves of pair instability supernovae in their restframe
Compared with a typical SN Ia (red SN 2001el), a Type Iip(blue. SN 1999em) and the hypernova SN 2006gy (green)
Red-shifted light curve of a bright pair-instability SN
Pulsational PairInstability Supernovae
Pulsations
Woosley et al. (2007; Nature 450, 390)
238 million light years away
Smith et al. (2007; ApJ 666, 1116)
Smith et al. (2007; ApJ 666, 1116)
Onset of instability
Woosley et al. (2007; Nature 450, 390)
At end of first pulse
Woosley et al. (2007; Nature 450, 390)
After 2nd pulse
Woosley et al. (2007; Nature 450, 390)
At final point
Woosley et al. (2007; Nature 450, 390)
Velocity and enclosed mass after second mass ejection - 110 solar mass model (74.6 at explosion)
Shock heating
Woosley et al. (2007; Nature 450, 390)
Light curves of the two outbursts (110 solar mass model)
Woosley et al. (2007; Nature 450, 390)
20061999 2012
Absolute R-band magnitudesof the 110 solar mass modelcompared with obsevations of “hypernova” SN 2006gy.
Instabilities will smooththese 1 D calculations.The brighter curve assumedtwice the velocity for all ejecta. (7.2 x 1050 ergbecomes 2.9 x 1051 erg)
Woosley et al. (2007; Nature 450, 390)
