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Supernova Light S.I.Blinnikov [email protected] IPMU, ITEP, SAI IPMU-Prsp09f – p.

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  • Supernova LightS.I.Blinnikov

    [email protected]

    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

    and

    Ekin = Ethermal =3

    2

    PdV

    For ER:

    P =1

    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

    a

    )1/3

    , P =1

    3aT 4 =

    a

    3

    (

    3

    4

    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

    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