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Supernova LightS.I.Blinnikov
IPMU, ITEP, SAI
IPMU-Prsp09f – p.
IPMU Colloquium18 Feb 2009
S.I.Blinnikov 1,2,3
1Institute for Theoretical andExperimental Physics (ITEP),Moscow
2IPMU, Tokyo
3Sternberg AstronomicalInstitute (SAI), Moscow
IPMU-Prsp09f – p.
SN1994D type Ia in NGC4526
IPMU-Prsp09f – p.
SNR Tycho in X-rays (Chandra)
IPMU-Prsp09f – p.
SN 2006gy
Ofek et al.2007, ApJL,
astro-ph/0612408)
Smith et al.2007, Sep. 10
ApJ, astro-ph/0612617)
IPMU-Prsp09f – p.
Brightest. Supernova. Everby N.Smith
IPMU-Prsp09f – p.
Most Luminous SN ever
IPMU-Prsp09f – p.
Luminous SN: too many photons?Now we know a few other SNe with peak luminosity evenhigher than SN 2006gy.
Total light 1051 ergs: 2 orders ofmag higher than normal corecollapsing SN and 1 order morethan brightest thermonuclear SNTo explain this light we inevitably involve large stellarmasses.I’ll try to explain why the evolution of stars with M > 10M⊙ isquite different from low mass stars, and what happens atM ∼ 100M⊙
IPMU-Prsp09f – p.
IPMU-Prsp09f – p.
IPMU-Prsp09f – p. 10
SN light productionMost powerful supernovae (SNe): what is the problem?
IPMU-Prsp09f – p. 11
SN light productionMost powerful supernovae (SNe): what is the problem?
Basics of stellar evolution: virial theorem and danger ofrelying on Tvir and vvir
IPMU-Prsp09f – p. 11
SN light productionMost powerful supernovae (SNe): what is the problem?
Basics of stellar evolution: virial theorem and danger ofrelying on Tvir and vvirPair-instability SNe and nonsense on them in Wikipedia
IPMU-Prsp09f – p. 11
SN light productionMost powerful supernovae (SNe): what is the problem?
Basics of stellar evolution: virial theorem and danger ofrelying on Tvir and vvirPair-instability SNe and nonsense on them in Wikipedia
Shock propagation and entropy creation inside a star
IPMU-Prsp09f – p. 11
SN light productionMost powerful supernovae (SNe): what is the problem?
Basics of stellar evolution: virial theorem and danger ofrelying on Tvir and vvirPair-instability SNe and nonsense on them in Wikipedia
Shock propagation and entropy creation inside a star
Radiation-dominated shocks and tenacious myths onX-rays from SN shocks
IPMU-Prsp09f – p. 11
Stellar evolutionHR (L− Teff) diagram needed for comparison with
observations
IPMU-Prsp09f – p. 12
Compression in centereven if Rout grows
1974ARA&A..12..215I
IPMU-Prsp09f – p. 13
Virial theoremFrom Baz’, Zeldovich, Perelomov...
δEtot ≡ δE = δ〈H〉 = 0in the first order of perturbation δψ – the variational principlein quantum case.Let us take a perturbation of the form:
ψ + δψ = α3N/2ψ({α~ri}), i = 1, . . . , N,
so the wave function is uniformly changed for all 3N spacecoordinates. The coefficient α3N/2 is from normalization tounity. Note, that α > 1 here corresponds to the compressionof the whole system.
IPMU-Prsp09f – p. 14
We have〈H〉 = 〈Ekin〉 + 〈U〉.
Now〈Ekin〉 → α2〈Ekin〉
because for non-relativistic (NR) particles, Ekin ∝ p2 ∝ 1/λ2The variation of U depends on the law of interaction. ForCoulomb (and Newton!) interactions
〈U〉 ∝∫
∑
i 6=k
ψ∗1
rikψ dN~r → α〈U〉.
IPMU-Prsp09f – p. 15
Thus〈H〉 → α2〈Ekin〉 + α〈U〉,
and variation of this gives
δ〈H〉 = 2αδα〈Ekin〉 + δα〈U〉 = 0,
so with α = 1 for unperturbed state we find
2〈Ekin〉 + 〈U〉 = 0.
This is the virial theorem for atomic Coulomb potential (andfor globular stellar clusters as well! See the use ofSchrödinger Eq. for stellar dynamics in Widrow & Kaser,1993)
IPMU-Prsp09f – p. 16
For all those systems (NR atoms or plasma, NR stars andstellar clusters)
E = 〈Ekin〉 + 〈U〉 = −〈Ekin〉
so the loss of total energy E corresponds to the growth ofthe kinetic energy 〈Ekin〉. The same is true for the internalenergy of matter if it is in the form of kinetic energy ofparticles.One can do U ∝ rk, but more important for us is extremelyrelativistic (ER) case: Ekin ∝ p ∝ 1/λ, then
〈Ekin〉 + 〈U〉 = 0.
IPMU-Prsp09f – p. 17
PressureFor NR:
P =2
3ε
and
Ekin = Ethermal =3
2
∫
PdV
For ER:
P =1
3ε
and
Ekin = Ethermal = 3
∫
PdV
and always
3
∫
PdV + 〈U〉 = 0IPMU-Prsp09f – p. 18
Central PressureNow, for U ∼ −GNM2/R, omitting all coefficients of orderunity, pressure and density in the center are:
Pc ≃GNM
2
R4, ρc ≃
M
R3.
and we findPc ≃ GNM
2
3ρ4/3c .
IPMU-Prsp09f – p. 19
Tc ∝M 2/3ρ1/3c in ND starsSo if we have a classical ideal plasma with P = RρT/µ,where R is the universal gas constant, and µ – meanmolecular mass,
Tc ≃GNM
2/3ρ1/3c µ
R .
With µ ≃ 1 for H-He fully ionized plasma we get for the SunTc ≃ 107 K ≃ 1 keV. This is OK for “virial” velocityv2vir ∼ ϕ ∼ rg/R⊙ ∼ 10−6 (although
√2 less than for
neutrals!). May be called “virial” Tvir for ions, but not forelectrons.
IPMU-Prsp09f – p. 20
Now check: Tc ∝M 2/3ρ1/3c
IPMU-Prsp09f – p. 21
Check: Tc ∝M 2/3ρ1/3c
IPMU-Prsp09f – p. 22
Not so for lower masses
IPMU-Prsp09f – p. 23
Degeneracy of electrons
IPMU-Prsp09f – p. 24
Degeneracy of electrons
IPMU-Prsp09f – p. 25
M > 10M⊙ never degenerate
IPMU-Prsp09f – p. 26
Compare with old Iben’s results1974ARA&A..12..215I
IPMU-Prsp09f – p. 27
Check: Tc ∝M 2/3ρ1/3c
IPMU-Prsp09f – p. 28
Check: Tc ∝M 2/3ρ1/3c
IPMU-Prsp09f – p. 29
If radiation dominates in P
When plasma isradiation-dominated (for massivestars), then, P ∝ T 4, and
Tc ∝M 1/6ρ1/3c .
IPMU-Prsp09f – p. 30
HR and Tc − ρc evolution
IPMU-Prsp09f – p. 31
Compute stars yourselfComputational Astrophysics:
http://rainman.astro.uiuc.edu/ddr/
The Digital Demo Room10000 stars evolve together – find on this site
7 stars of masses 20M⊙ < M < 80 evolve in a combined runand explode as SNe – produced on this site – click here
IPMU-Prsp09f – p. 32
http://rainman.astro.uiuc.edu/ddr/
The carbon-oxygen cores of low mass stars turn out to bedegenerate at the moment when the carbon burningbegins. The temperature of their interiors is also stronglyaffected by the neutrino energy losses. Should the carbonburning only begin in degenerate conditions, it acquires aviolent, explosive nature giving rise to the explosion of TypeIa supernovae.
IPMU-Prsp09f – p. 33
On hydrodynamical instability
Equilibrium requires (in Newtonian gravity):
Pc ≃ GNM 2/3ρ4/3c .
This implies that adiabatic exponent
γ < 4/3 may lead to a hydrodynamic
instability.
IPMU-Prsp09f – p. 34
Mechanical stability
IPMU-Prsp09f – p. 35
Relativistic particles
lead to γ → 4/3We have γ ∼ 4/3 due to high entropy S (photons and e+e−pairs).At low S → 0 we have γ → 4/3 due to high Fermi energy ofdegenerate electrons at high density ρ.
IPMU-Prsp09f – p. 36
Causes for a collapse: pairsFor very massive stars the radiation pressure aT 4/3 mustbe much larger than RρT .Here per gram
Eth = aT4/ρ
and fromTS = Eth + P/ρ for µ = 0,
we find per unit mass
S =4
3
aT 3
ρ.
IPMU-Prsp09f – p. 37
Photons and ...
T =
(
3
4
Sρ
a
)1/3
, P =1
3aT 4 =
a
3
(
3
4
Sρ
a
)4/3
,
i.e. P ∝ ρ4/3 for constant S, and γ = 4/3. When T >∼ 0.1mec2for small µ in non-degenerate gas the pairs (e+e−) are bornintensively, so for T ≫ mec2 the total thermal energy
Ethρ = aT4 +
7
4aT 4 =
11
4aT 4,
(7a/8) is added per each polarization of fermions.Exact formulae see, e.g., SB,Dunina-Barkovskaya,DKN,1996, ApJS.
IPMU-Prsp09f – p. 38
. . . and e+e− pairspressure
P =11
12aT 4,
and entropy per gram
S =11
3
aT 3
ρ.
Thus for T ≫ mec2 again P ∼ ρ4/3, but the coefficient issmaller
P =11
12aT 4 =
11a
12
(
3
11
Sρ
a
)4/3
,
so in between the slope logP – log ρ must be less than 4/3.
IPMU-Prsp09f – p. 39
Pair instability
A radiation dom-inated star wasalready at the vergeof the loss of the sta-bility (P ∝ ρ4/3), andnow it is unstable if(γ < 4/3).
IPMU-Prsp09f – p. 40
Higher mass means higher Tcfor the same ρ, hence pair creation
IPMU-Prsp09f – p. 41
Open evolution codePrevious plot is taken from here
Paxton: P.Eggleton evolution codeCentre Temperature-Density Tracks for Metallicity Z = 0.02.
The “He” symbols show where the net of power fromnuclear reactions beyond hydrogen burning minus neutrino
losses from all sources reaches the break-even point.
IPMU-Prsp09f – p. 42
http://theory.kitp.ucsb.edu/~paxton/
Adiabatic γ for pairsGary S. Fraley 1968. Pair-instability SNe
IPMU-Prsp09f – p. 43
Adiabatic γ for pairsD.K.Nadyozhin 1974, see SB,Dunina-Barkovskaya,DKN,
1996, ApJS
IPMU-Prsp09f – p. 44
Umeda and Nomoto 2007
IPMU-Prsp09f – p. 45
Woosley et al. s103
This gives the Most Luminous Supernovae!IPMU-Prsp09f – p. 46
SN light production in shocksRadiation-dominated shocks and tenacious myths onX-rays from SN shocks
IPMU-Prsp09f – p. 47
SN light production in shocksRadiation-dominated shocks and tenacious myths onX-rays from SN shocks
Shock propagation and entropy creation inside a star
IPMU-Prsp09f – p. 47
SN light production in shocksRadiation-dominated shocks and tenacious myths onX-rays from SN shocks
Shock propagation and entropy creation inside a star
Shock breakout, if any
IPMU-Prsp09f – p. 47
SN light production in shocksRadiation-dominated shocks and tenacious myths onX-rays from SN shocks
Shock propagation and entropy creation inside a star
Shock breakout, if any
Diffusion of photons and cooling of ejecta
IPMU-Prsp09f – p. 47
Models for powerful light
IPMU-Prsp09f – p. 48
Models for powerful lightRadiative shock, Chugai, Blinnikovea’04, Woosley ea’07 ...
(photosphere)
IPMU-Prsp09f – p. 48
Models for powerful lightDense shell Chugai, Blinnikov ea’04
IPMU-Prsp09f – p. 48
Models for powerful lightDense shell Chugai, Blinnikov ea’04 Radioactive, Nomoto, Tominaga ea’07
-22
-21
-20
-19
-18
-17 0 50 100 150 200 250 300
Mag
nitu
des
Days
166M E65 Ni1506gy r Ofek
06gy R Smith53M E64 Ni15
IPMU-Prsp09f – p. 48
Models for powerful lightDense shell Chugai, Blinnikov ea’04 Radioactive, Nomoto, Tominaga ea’07
-22
-21
-20
-19
-18
-17 0 50 100 150 200 250 300
Mag
nitu
des
Days
166M E65 Ni1506gy r Ofek
06gy R Smith53M E64 Ni15
Diffusion, McCray, Smith’07SN 2006gy
Type II
IPMU-Prsp09f – p. 48
Models for powerful lightDense shell Chugai, Blinnikov ea’04 Radioactive, Nomoto, Tominaga ea’07
-22
-21
-20
-19
-18
-17 0 50 100 150 200 250 300
Mag
nitu
des
Days
166M E65 Ni1506gy r Ofek
06gy R Smith53M E64 Ni15
Diffusion, McCray, Smith’07SN 2006gy
Type II
Atmospheric, Dessart ea’08
IPMU-Prsp09f – p. 48
SN 2006gy
Ofek et al.2007, ApJL,
astro-ph/0612408)
Smith et al.2007, Sep. 10
ApJ, astro-ph/0612617)
IPMU-Prsp09f – p. 49
Smith et al. SN06gy spectraNarrow lines: SNIIn
IPMU-Prsp09f – p. 50
Another LC set with SN IIn
IPMU-Prsp09f – p. 51
Problems posed by SN2006gy
Total light 1051 ergs: 2 orders of mag higher than
normal SN II
If the source is radioactive material, then huge
amount, high explosion energy
If the source is a shock wave, why X-ray was
weak
What evolution leads to any of those outcomes?
IPMU-Prsp09f – p. 52
Spectra of SNe IIn
IPMU-Prsp09f – p. 53
SN light
Many photons ⇒ high entropy: S ∼ nγ/nb.Thus a source of S is needed for luminous
SNe: either radioactivity or shocks.
IPMU-Prsp09f – p. 54
Chaos ⇒ Order in expansion: 0
IPMU-Prsp09f – p. 55
Chaos ⇒ Order in expansion: 3
IPMU-Prsp09f – p. 56
Chaos ⇒ Order in expansion: 10
IPMU-Prsp09f – p. 57
Chaos ⇒ Order in expansion: 30
IPMU-Prsp09f – p. 58
Chaos ⇒ Order in expansion: 100
IPMU-Prsp09f – p. 59
Sources of photons in SN I
radioactivedecays 56Ni→56Co (followingNadyozhin)
IPMU-Prsp09f – p. 60
Sources of photons in SN I
radioactivedecays 56Co→56Fe
IPMU-Prsp09f – p. 61
Sources of entropy in SN II: ShocksShock inside the star remains in adiabatic phase while opticaldepth,
δR
l>
c
D,
where δR is the distance from the shock to the photosphere(Ohyama 1963, also Imshennik V.S., Morozov Yu.I. 1964)When
δR
l
Shock Luminosity in SN II
Shocks:Differentradii atshock-
breakoutepoch
IPMU-Prsp09f – p. 63
Shock lg T in SN II
IPMU-Prsp09f – p. 64
Pure diffusion models for SNIIn
The entropy produced byshock in a cloud may be usedas a reservoir of photons.If the cloud is large in radiusand in mass, then highluminosity may be achieved
IPMU-Prsp09f – p. 65
N.Smith: SN06gy data
Model byN.Smith &R.McCrayis shown
SN 2006gy
Type II
IPMU-Prsp09f – p. 66
SN II LC TheoryImshennik, Nadyozhin & Grasberg (1964-1971)
IPMU-Prsp09f – p. 67
Huge PreSN LC
IPMU-Prsp09f – p. 68
Blinnikov & Popov 1993advancing Arnett’s analytical theory
IPMU-Prsp09f – p. 69
Blinnikov & Popov 1993
IPMU-Prsp09f – p. 70
Light Curve Smith McCray 1
IPMU-Prsp09f – p. 71
Photoph. speed SM 1
IPMU-Prsp09f – p. 72
Light Curve Smith McCray 7
IPMU-Prsp09f – p. 73
Photoph. speed SM 7
IPMU-Prsp09f – p. 74
Type IIP photosphere
almost at rest -not much
expanding in R
Recombinationfront movinginside in Mr
IPMU-Prsp09f – p. 75
‘Visible’ disk of SN IIP
IPMU-Prsp09f – p. 76
Difficulties of diffusion models
High luminosity leads to a coolingwave – the evolution of light andphotospheric velocity is too fast forSN 2006gy.If the cloud is large in radius and in mass,a strong shock is radiative anddegenerates into a thin shell: initial heatingmust be done by a mysterious source ofheating uniformly distributed throughoutthe cloud
IPMU-Prsp09f – p. 77
SN IIn: the brightest light
IPMU-Prsp09f – p. 78
SN 1994W line power
IPMU-Prsp09f – p. 79
SN1994W & SN2006gy structure
(photosphere)
IPMU-Prsp09f – p. 80
SN 1994W line profile models
IPMU-Prsp09f – p. 81
Shocks in SNe IIn
A long liv-ing shock:an exampleof SN1994wof type IIn.Density as afunction of theradius r in twomodels at day30. The struc-ture tends toan isothermalshock wave.
IPMU-Prsp09f – p. 82
Light Curve SN 1994W
Light curvesfor the runsn94w58.Fluxes in BVfilters converge- contraryto SN II, agood featureto distinguishSN IIn beforespectroscopicobservations.
IPMU-Prsp09f – p. 83
Spectrum modeling R(t)
Chugai’s kine-matic modelThe CS density(top panel)and evolutionof the radiusand velocityof the cooldense shell(two bottompanels)
IPMU-Prsp09f – p. 84
On the same plot
Broken line is from spectra, solid – from hydro LC model.
IPMU-Prsp09f – p. 85
Light Curve SN 1994W
IPMU-Prsp09f – p. 86
Light Curve b SN 1994W
IPMU-Prsp09f – p. 87
‘Visible’ disk of SN IIn
IPMU-Prsp09f – p. 88
Observed spectrum of SN 1994W
IPMU-Prsp09f – p. 89
SN Light Source
Cooling
no source of S
type II
Permanent heatingby shocktype IIn
Fading heatingby radioactivity
type Ia, Ib/c
IPMU-Prsp09f – p. 90
SN Light Source
Cooling
no source of S
type II
Permanent heatingby shocktype IIn
Fading heatingby radioactivity
type Ia, Ib/c
IPMU-Prsp09f – p. 91
SN Light Source
Cooling
no source of S
type II
Permanent heatingby shocktype IIn
Fading heatingby radioactivity
type Ia, Ib/c
IPMU-Prsp09f – p. 92
‘Visible’ of SN IIP c
IPMU-Prsp09f – p. 93
Baklanov et al. (2005)SN 1999em D= 12.38 Mpc
IPMU-Prsp09f – p. 94
Type IIP too weak for SN06gy
A huge 1000R⊙ RSG produces a rather weak SN II. Is itfeasible for radioactivity? Yes, but for tens of M⊙ in 56Ni.One has to go to VERY massive stars.
IPMU-Prsp09f – p. 95
Very Massive Stars
The bigger they comeThe harder they fall,One and all
Jimmy Cliff (1973)as recorded by Willy Nelson
The lives and deaths of the first stars
Alex Heger (LANL, UCSC)S. E. Woosley (UCSC)Weiqun Zhang (Stanford)Candace Church (UCSC)Sergei Blinnikov (ITEP)
IPMU-Prsp09f – p. 96
Four kinds of deaths
With some uncertainty about exact demarcations, one can delineate four kinds of deaths for non-rotating helium stars.(For rotation decrease main sequence mass 10 - 20%)
He Core Main Seq. Mass Supernova Mechanism
2 d M d40 10d M d95 Fe core collapse to neutron star
or a black hole
40d M d60 95d M d 130 Pulsational pair instability followed
by Fe core collapse
60d M d137 130d M d260 Pair instability supernova
M t137 M t260 Black hole. Possible GRB
IPMU-Prsp09f – p. 97
Very Massive Stars
approximate
Eexpl56Ni
Fe-richFe-poor
Heger and Woosley (2002)
IPMU-Prsp09f – p. 98
Bolometric for ed250
�
Re = 1.3 u 1014 cm; M
He = 127 Me
M(56Ni) = 38.9 Me
KEf
= 86 u 1051 erg
Calculations by Sergei Blinnikov
Brighter still are the biggerexplosionsthat make lotsof 56Ni
IPMU-Prsp09f – p. 99
UBVR for ed250 and SN06gy
IPMU-Prsp09f – p. 100
K. Nomoto, N. Tominaga,
M. Tanaka,K. Maeda,H. Umeda, 2007
-22
-21
-20
-19
-18
-17 0 50 100 150 200 250 300
Mag
nitu
des
Days
166M E65 Ni1506gy r Ofek
06gy R Smith53M E64 Ni15
IPMU-Prsp09f – p. 101
Pair instability
A radiation dom-inated star wasalready at the vergeof the loss of the sta-bility (P ∝ ρ4/3), andnow it is unstable if(γ < 4/3).
IPMU-Prsp09f – p. 102
Umeda and Nomoto 2007
IPMU-Prsp09f – p. 103
Woosley et al. s103
This gives the Most Luminous Supernovae!IPMU-Prsp09f – p. 104
SN-repeaters
IPMU-Prsp09f – p. 105
SN-repeaters2
IPMU-Prsp09f – p. 106
SN-repeaters3
IPMU-Prsp09f – p. 107
Stella: LCs for SN2006gy
IPMU-Prsp09f – p. 108
Double explosion: old ideaGrasberg & Nadyozhin (1986)
IPMU-Prsp09f – p. 109
Hydro structure 90 d
IPMU-Prsp09f – p. 110
‘Visible’ disk of SN 2006gy
IPMU-Prsp09f – p. 111
‘Visible’ disk of SN 2006gy c
IPMU-Prsp09f – p. 112
90 d, mass coordinate
IPMU-Prsp09f – p. 113
Velocity at τ = 0.1
IPMU-Prsp09f – p. 114
LCs for doubled ρ
IPMU-Prsp09f – p. 115
SN 2006tf and models
IPMU-Prsp09f – p. 116
LCs for doubled velocity
IPMU-Prsp09f – p. 117
Higher res., tail observed
K. Kawabata,et al. (2008),ApJ subm.,CourtesyM.Tanaka
IPMU-Prsp09f – p. 118
Soft X-ray from SN06gy?
IPMU-Prsp09f – p. 119
Postshock temperatureT2 behind a strong shock, constant γ, small losses/gains Q,for ideal gas, neglecting radiation, is
T2 =2(γ − 1)u21µ(γ + 1)2R =
3u21µ
16R for γ = 5/3 .
If we put here D8 = u1/108cm/s, then D8 is the shock speedin thousand km/s and we get
T2 = 2.25 × 107µD28in K or
T2(keV) = 1.94µD28in keV. Here µ = A/(1 + Z) for plasma (sincen = nbaryon/µ = nionA/µ = nion + ne = nion + Znion).
IPMU-Prsp09f – p. 120
Hydro structure 120 d
X-rays from the shock cannot go out yet, the matter is too dense.IPMU-Prsp09f – p. 121
120 d, mass coordinate
IPMU-Prsp09f – p. 122
Other X-ray SNe IIn
SN galaxy dist. discoveryMpc day
1986J NGC 891 9.6 3,3001988Z +03-28-022 89 2,3701994W NGC 4041 25 1,1801995N -2-38-017 24 4401998S NGC 3877 17 6782005ip NGC 2096 30 4902005kd PGC14370 64 450
All SN IIn discovered late in X-ray!
IPMU-Prsp09f – p. 123
‘Visible’ disk of SN 2006gy
IPMU-Prsp09f – p. 124
‘Visible’ disk of SN 2006gy c
IPMU-Prsp09f – p. 125
MC probable d to SN 2006gy
IPMU-Prsp09f – p. 126
Multi-D is a mustfor next steps in theoretical modeling
IPMU-Prsp09f – p. 127
ConclusionsRadiating shocks are most probable sources of light inmost luminous supernovae like SN2006gy.
The medium for the shining shock to propagate isnaturally produced in massive star evolution due toviolent non-linear pulsations when e−e+ pairs becomeappreciable in the pressure in their interiors.
The supercritical radiative shock may be well belowX-ray T for a long time.
Most luminous SN 2006gy events may be observed athigh z [for years due to (1 + z)] and may be useful asdirect distance indicators in cosmology.
IPMU-Prsp09f – p. 128
IPMU ColloquiumSN1994D type Ia in NGC4526SNR Tycho in X-rays (Chandra)SN~2006gyBrightest. Supernova. EverMost Luminous SN everLuminous SN: too many photons?SN light productionSN light productionSN light productionSN light productionSN light production
Stellar evolutionCompression in centerVirial theoremPressureCentral Pressure$T_c propto M^{2/3}ho _c^{1/3}$ in ND starsNow check: $T_c propto M^{2/3}ho _c^{1/3}$Check: $T_c propto M^{2/3}ho _c^{1/3}$Not so for lower massesDegeneracy of electronsDegeneracy of electrons$M> 10 msun $ never degenerateCompare with old Iben's resultsCheck: $T_c propto M^{2/3}ho _c^{1/3}$Check: $T_c propto M^{2/3}ho _c^{1/3}$If radiation dominates in $P$HR and $T_c - ho _c$ evolutionCompute stars yourselfOn hydrodynamical instabilityMechanical stabilityRelativistic particlesCauses for a collapse: pairsPhotons and ... $ldots $ and $e^+e^-$ pairsPair instabilityHigher mass means higher $T_c $Open evolution codeAdiabatic $gamma $ for pairsAdiabatic $gamma $ for pairsUmeda and Nomoto 2007Woosley et al. s103SN light production in shocksSN light production in shocksSN light production in shocksSN light production in shocks
Models for powerful lightModels for powerful lightModels for powerful lightModels for powerful lightModels for powerful lightModels for powerful light
SN~2006gySmith et al. SN06gy spectraAnother LC set with SN~IInProblems posed by SN2006gySpectra of SNe~IInSN lightChaos $Rightarrow $ Order in expansion: 0Chaos $Rightarrow $ Order in expansion: 3Chaos $Rightarrow $ Order in expansion: 10Chaos $Rightarrow $ Order in expansion: 30Chaos $Rightarrow $ Order in expansion: 100Sources of photons in SN~ISources of photons in SN~ISources of entropy in SN~II: ShocksShock Luminosity in SN~IIShock $lg T$ in SN~IIPure diffusion models for SNIInN.Smith: SN06gy dataSN~II LC TheoryHuge PreSN LCBlinnikov & Popov 1993Blinnikov & Popov 1993Light Curve Smith McCray 1Photoph. speed SM 1Light Curve Smith McCray 7Photoph. speed SM 7Type IIP photosphere`Visible' disk of SN~IIPDifficulties of diffusion modelsSN~IIn: the brightest lightSN~1994W line powerSN1994W & SN2006gy structureSN~1994W line profile modelsShocks in SNe~IInLight Curve SN~1994WSpectrum modeling $R(t)$On the same plotLight Curve SN~1994WLight Curve b SN~1994W`Visible' disk of SN~IInObserved spectrum of SN~1994WSN Light SourceSN Light SourceSN Light Source`Visible' of SN~IIP cBaklanov et al.~(2005)Type IIP too weak for SN06gyVery Massive StarsFour kinds of deathsVery Massive StarsBolometric for ed250UBVR for ed250 and SN06gyK.~Nomoto, N.~Tominaga, Pair instabilityUmeda and Nomoto 2007Woosley et al. s103SN-repeatersSN-repeaters2SN-repeaters3Stella: LCs for SN2006gyDouble explosion: old ideaHydro structure 90 d`Visible' disk of SN~2006gy`Visible' disk of SN~2006gy c90 d, mass coordinateVelocity at $au =0.1$LCs for doubled $ho $ SN~2006tf and modelsLCs for doubled velocityHigher res., tail observedSoft X-ray from SN06gy?Postshock temperatureHydro structure 120 d120 d, mass coordinateOther X-ray SNe~IIn`Visible' disk of SN~2006gy`Visible' disk of SN~2006gy cMC probable $d$ to SN~2006gyMulti-D is a must Conclusions