Äîêóìåíò âçÿò èç êýøà ïîèñêîâîé ìàøèíû. Àäðåñ îðèãèíàëüíîãî äîêóìåíòà : http://nuclphys.sinp.msu.ru/nseminar/15.05.12.pdf Äàòà èçìåíåíèÿ: Thu May 17 13:47:37 2012 Äàòà èíäåêñèðîâàíèÿ: Tue Oct 2 00:18:11 2012 Êîäèðîâêà: Ïîèñêîâûå ñëîâà: emission nebula

Pair-instability
..
sergei.blinnikov@itep.ru

ITEP, SAI, also partly IPMU

15y12-Prosper ­ p. 1


SINP MSU, 15th May 2012
S.I.Blinnikov
1 1,2,3

Institute for Theoretical and Experimental Physics (ITEP), Moscow
2

Sternberg Astronomical Institute (SAI), MSU, Moscow

3

work partly done at IPMU, Tokyo University, Kashiwa
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Supernova SN1994D in NGC4526
Shocks are not important for light in "Nobel prize" SNe Ia

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SN 2006gy
Ofek et al. 2007, ApJL Smith et al. 2007, ApJ

Shocks are vital for explaining light of those superluminous events for many months...

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SNR Tycho in X-rays (Chandra)

...and thousands of years in SNRs
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Supernova: order of events
Core collapse (CC) or explosion

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Supernova: order of events
Core collapse (CC) or explosion Neutrino/GW signal, accompanying signals

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Supernova: order of events
Core collapse (CC) or explosion Neutrino/GW signal, accompanying signals Shock creation if any, propagation and entropy production inside a star

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Supernova: order of events
Core collapse (CC) or explosion Neutrino/GW signal, accompanying signals Shock creation if any, propagation and entropy production inside a star Shock breakout (!)

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Supernova: order of events
Core collapse (CC) or explosion Neutrino/GW signal, accompanying signals Shock creation if any, propagation and entropy production inside a star Shock breakout (!) Diffusion of photons and cooling of ejecta

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Core-Collapse-SN (CCSN)
Standard description of Chronology 1 sec: Core collapse, bounce, or SASI) , or rotMHD, shock revival 1 min to 1 day: shock propagates and breaks out (1st EM signature). Fallback? NS vs. BH formation? Mins to days: Final ejecta acceleration to homology (velocity u r )
)

Standing accretion shock instability

Actually some weak EM signals are inevitably produced before shock breakout

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Burning in center and in shells

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Many shells

next few slides from Raffelt (2010) and other sources

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Newborn NS

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NS energy estimates

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First messengers of explosions
Neutrino?

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First messengers of explosions
Neutrino? Gravitational waves?

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First messengers of explosions
Neutrino? Gravitational waves? Radio waves? At least in atmospheric explosions

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First messengers of explosions
Neutrino? Gravitational waves? Radio waves? At least in atmospheric explosions Shock breakout

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SN classification
thermonuclear core collapse

I SiII
yes no

no

H
yes

yes

II
IIL

IIP
curve light shape

HeI
no

IIb

Ia

Ib Ic

Ib/c pec IIn
hypernovae

strong ejecta-CSM interaction

Turrato 2003
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Extremely bright Type IIn SNe

V-band (Drake et al. 2010)

SN1987A and a typical SNII below the frame!

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H-poor superluminous SNe
Quimby et al. 2011
Absolute u-band AB magnitude -22
SCP 06F6 SN 2005ap PTF09cnd PTF09cwl PTF09atu PTF10cwr

-20

-18

-16
S SS S

-50

0 50 100 Rest-frame phase (days)

150

Still enigmatic. Most probably explained by a long living radiative shock. No better model is suggested
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Supernova 1987A Neutrinos

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SN 1987A Neutrinos
Ten neutrino events were detected in a deep mine neutrino detection facility in Japan which coincided with the observation of Supernova 1987A. They were detected within a time interval of about 15 seconds against a background of lower energy neutrino events. A similar facility, IMB in Ohio detected 8 neutrino events in 6 seconds. These observations were made

18 hours

before the first optical sighting of the supernova.

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Superlumnal neutrino cartoons...

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Longo PRD 36(1987)3276

Distance = 1.6 â 105 ly, t 3h , hence
|(c - c )/c| 3h /(1.6 â 105 â 365 â 24) = 2 â 10-
9

Where does t 3h come from? Could the constraint be improved?

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SN1987A discovery
Timing (times in Universal Time) 7:36, 23 February, neutrinos observed 9:30, 23 February Albert Jones, amateur astronomer, observes Tarantula Nebula in LMC He sees nothing unusual 10:30, 23 February Robert McNaught photographs LMC When plate is developed, SN1987A is there. Some 20 hours later, Ian Shelton's discovery.

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SN87A early observations
Blinnikov with K.Nomoto ea

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SN87A early observations

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Improvement of c constraint
If the flash at shock breakout were observed we would get
|(c - c )/c| 2 â 10-
10

Much better improvement is possible in principle! If a precollapse suspect is monitored and its prompt quake is registered e.g. in radio simultaneously with and/or GW signal.

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detectors

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Next generation detectors

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Gravitational Waves from CCSNe
http://numrel.aei.mpg.de/images

These images are copyright of AEI, ZIB, LSU and SISSA

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GW detectors

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detectors

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GW LIGO estimates

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SN 2006gy
Ofek et al. 2007, ApJL, astroph/0612408) Smith et al. 2007, Sep. 10 ApJ, astroph/0612617)

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Brightest. Supernova. Ever
by N.Smith

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It was Most Luminous SN ever

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Extremely bright Type IIn SNe

V-band (Drake et al. 2010)

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Luminous SN: too many photons?
Now we know a few other SNe with peak luminosity even higher than SN 2006gy.

Total light 10 ergs: 2 orders of mag higher than normal core collapsing SN and 1 order more than brightest thermonuclear SN
To explain this light we inevitably involve large stellar ma s s e s . I will try to explain why the evolution of stars with M > 10M is quite different from low mass stars, and what happens at M 100M

51

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Stellar evolution
HR (L - Teff ) diagram needed for comparison with observations

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CompressionniinR cenwer ts eve f gro
out
1974ARA&A..12..215I

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Central Pressure
Omitting all coefficients of order unity, pressure and density in the center are:
M GN M 2 , c 3 . Pc 4 R R

and we find
Pc GN M
2 3

4/3 c

.

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Tc M

2/3 1/3 c

in ND stars

So if we have a classical ideal plasma with P = RT /µ, where R is the universal gas constant, and µ ­ mean molecular mass,
Tc GN M
2/3 1/3 c

µ

R

.

With µ 1 for H-He fully ionized plasma we get for the Sun Tc 107 K 1 keV.

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Now check: Tc M

2/3 1/3 c

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Check: Tc M

2/3 1/3 c

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Not so for lower masses

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Degeneracy of electrons

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Degeneracy of electrons

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M > 10M never degenerate

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Compare with old Iben's results
1974ARA&A..12..215I

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Check: Tc M

2/3 1/3 c

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Check: Tc M

2/3 1/3 c

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If radiation dominates in P When plasma is radiation-dominated (for massive 4 stars), then, P T , and Tc M
1/6 1/3 c

.

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HR and Tc - c evolution

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Compute stars yourself
Computational Astrophysics: http://rainman.astro.uiuc.edu/ddr/ The Digital Demo Room 10000 stars evolve together ­ find on this site ­ click here 7 stars of masses 20M < M < 80 evolve in a combined run and explode as SNe ­ find on this site ­ click here

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The carbon-oxygen cores of low mass stars turn out to be degenerate at the moment when the carbon burning begins. The temperature of their interiors is also strongly affected by the neutrino energy losses. Should the carbon burning only begin in degenerate conditions, it acquires a violent, explosive nature giving rise to the explosion of Type Ia supernovae.

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On hydrodynamical instability
Equilibrium requires (in Newtonian gravity): Pc GN M
2/3 4/3 c

.

This implies that adiabatic exponent < 4/3 may lead to a hydrodynamic instability.

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Mechanical stability
lg P
5/3

S1 > S2 > S
4/3

3

â

M = co n s t

lg

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Relativistic particles
lead to 4/3
We have 4/3 due to high entropy S (photons and e+ e- pairs). At low S 0 we have 4/3 due to high Fermi energy of degenerate electrons at high density .

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Causes for a collapse: pairs
For very massive stars the radiation pressure aT 4 /3 must be much larger than RT . Here per gram Eth = aT 4 / and from
T S = Eth + P /

for

µ = 0,

we find per unit mass
4 aT 3 S= . 3

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Photons and ...
T= 3 S 4a
1/3

,

14a P = aT = 3 3

3 S 4a

4/3

,

> i.e. P 4/3 for constant S , and = 4/3. When T 0.1me c2 for small µ in non-degenerate gas the pairs (e+ e- ) are born intensively, so for T me c2 the total thermal energy

7 4 11 4 Eth = aT + aT = aT , 4 4
4

(7a/8) is added per each polarization of fermions. Exact formulae see, e.g., SB,Dunina-Barkovskaya,DKN, 1996, ApJS.

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. . . and e e pairs
pressure
11 4 P = aT , 12

+-

and entropy per gram
11 aT 3 . S= 3

Thus for T me c2 again P 4/3 , but the coefficient is smaller 4/3 11 4 11a 3 S P = aT = , 12 12 11 a so in between the slope log P ­ log must be less than 4/3.

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Pair instability
A dom was the loss (P now ( < radiation inated star already at verge of the of the stability 4/3 ), and it is unstable if 4/3).

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Open evolution code
Hertzsprung-Russell and Center Temperature-Density Tracks for Metallicity Z = 0.02. The "He" symbols show where the net of power from nuclear reactions beyond hydrogen burning minus neutrino losses from all sources reaches the break-even point. Paxton: P.Eggleton evolution code

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Higher for thesame ,eenns ahcgationr T ma s s m h a ce p ir ire he

c

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Open evolution code
Previous plot is taken from here Paxton: P.Eggleton evolution code Centre Temperature-Density Tracks for Metallicity Z = 0.02. The "He" symbols show where the net of power from nuclear reactions beyond hydrogen burning minus neutrino losses from all sources reaches the break-even point. -- Roni Waldman arXiv:0806.3544. Better looking plots below are from Roni Waldman's eprint arXiv:0806.3544.

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Massive stars and their He-cores
10 9.9 9.8 9.7 9.6 9.5 9.4 log(Tc) 9.3 9.2 9.1 9 8.9 8.8 8.7 8.6 8.5 3 4 5 6 7 log(c) 8 M 20 He 8 M 80 He 36 9 10 Pair instability Fe disintegration

Each line is labeled "M" for stellar models and "He" for He-core models, followed by the mass of the model or of the core. Here are stars that reach core collapse avoiding pair instability.

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3 outcomes of pair-instability
10

Here are only He-core models, labeled by "He" and the mass of the core. They all reach pair instability, subsequently experiencing 1) pulsations (He48), 2) complete disruption (He80), or 3) direct collapse (He160).

9.9 9.8 9.7 9.6 9.5 9.4 log(Tc) 9.3 9.2 9.1 9 8.9 8.8 8.7 8.6 8.5 3 4 5 6 7 log(c) He 48 He 80 He 160 8 9 10 Pair instability Fe disintegration

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< 4/3 domain in T - plane
Gary S. Fraley 1968. Pair-instability SNe

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Adiabatic for pairs at very low density
D.K.Nadyozhin 1974, see SB,Dunina-Barkovskaya,DKN, 1996, ApJS

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Umeda and Nomoto 2007

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Woosley et al. 2007, 103 M star
This gives the Most Luminous Supernovae (!), because, instead of one SN explosion, we have several ma s s ejections and collisions of mass shells which produce bright radiating shocks.

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SN IIn structure, Chugai, SB ea'04

(photosphere)

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Shocks in SNe IIn
A long living shock: an example for SN1994w of type IIn. Density as a function of the radius r in two models at day 30. The structure tends to an isothermal shock wave.

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Woosley, Blinnikov, Heger, s103

Pulsational pair instability may give the Most Lumio5u1s-PSsurpe.r72ovae! n 1 y 2 ro pe ­ p n


Two mass ejections

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Light curve for SN2006gy
from Woosley, SB, Heger (2007)

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Stella: LCs for SN2006gy
new runs

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Double explosion: old idea
Grasberg & Nadyozhin (1986)

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Hydro structure 60 d

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60 d, mass coordinate

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`Visible' disk of SN 2006gy

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Star formation rate = SFR
Smartt S. J., 2009, ARAA, 47, 63
100 R(z) (number s )
-1

NS births ccSNe

80 60 40 20 0 0

2

4 6 redshift z

8

10

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Nearby candidate:
Betelgeuse in ORION ­ distance 130 pc

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Neutrino warning

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Neutrino emission

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From Odrzywolek et al.

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Neutrinos: Milky Way warning
Red circle is expected range for GADZOOKS!/Super-Kamiokande detector Green circle is expected range for Gd-loaded 0.5 Mt water detector (UNO, Hyper-Kamiokande, LAGUNA) Blue circle is expected range for hypothetical 10 Mt under water detector (TITAN-D, under water balloon) Yellow circle is expected range for futuristic "Gigaton Array" detector -- for three hour warning range is much larger than Galaxy radius

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Neutrinos: 1 day MW warning

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3 hours MW warning

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Summary
Radiating shocks are most probable sources of light in most luminous supernovae of type IIn like SN2006gy Most luminous SN IIn events may be observed at high z [for years due to (1 + z )] and may be useful as direct, primary, distance indicators in cosmology

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Conclusions-1
The shock wave which runs through rather dense matter surrounding an exploding star can produce enough light to explain very luminous SN events. No 56 Ni is needed in this case to explain the light cur ve near maximum light (some amount may be needed to explain light curve tails). We need the explosion energy of only 2-4 Bethe for the shell with M = 3 - 6M and R 1016 cm. NARROW LINES MAY NOT BE PRODUCED!

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Conclusions-2
Questions on the latest phases of star evolution arise: Is it possible to form so big and dense envelopes? And how? Time scale for such a formation How far can the envelope extend? Density and temperature profiles inside the envelope right before the explosion Question to observations: try to find traces of such shells for bright explosions. (There are spectral evidence of circumstellar shells for type IIn and Ibn SNe. Is it possible to find C­O envelopes as well?)

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Conclusions-3
Many technical problems in light curve calculations: line opacities; dimensionality: 3D is preferable, since the envelope can most probably be clumpy; NLTE spectra

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Acknowledgements
Our work is supported by Grants of the Government of the Russian Federation 11.G34.31.0047, RFBR 10-02-00249, 10-02-01398, RF Sci. School 3458.2010.2, 3899.2010.2, by a grant IZ73Z0-128180/1 of the Swiss National Science Foundation (SCOPES), and in Germany by MPA guest program. This research has been partly done at IPMU, Tokyo University, and supported in part by World Premier International Research Center Initiative, MEXT, Japan

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