SUPERNOVA NEUTRINOS AT ICARUS G. Mangano INFN, Napoli.
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SUPERNOVA NEUTRINOS AT
ICARUS
G. Mangano
INFN, Napoli
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Summary
- SN explosion dynamics- Neutrino spectra and overall features- SN 1987A at Kamiokande and IMB- SN & ICARUS- SNO, SK, LVD - Oscillations- Issues to be studied
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H-R DIAGRAM for M3
H burning
turn-off
growing He core
He core burningHe flash
He and H shell burning
white dwarfs
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SN explosion dynamicsProgenitor Proto Neutron Star
~ 109 g/cm3 ~ 3 1014 g/cm3
T ~ 1010 K T ~ 1011 K
MFe ~ 1.4 M MPNS ~ 1.4 – 1.7 M
RFe ~ 6 103 Km RPNS ~ 10 - 15 Km
Energetics
E ~ G MNS2/RNS =1.6 1053 erg (MNS/ M)2 (10 km/RNS)
99% neutrinos
1% kinetic energy
0.01% photons !!
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Evolved massive stars (M> 8 M) have a degenerate core of iron group elements (the most tightly bound nuclei) no further nuclear burning phase
at T125 MeV iron photodissociation: instability and collapse begins
Pressure lost via e- capture on nuclei
Inner core collapse is homologous (v/r 400-700 s-
1)
subsonic for the inner part
supersonic for the outer part
nFe 41356
eAZ
AZ XXe
1
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Neutrino sphere: diffusion time (neutral current interactions on nuclei) larger than collapse time:
’s are trapped in a degenerate sea (YL0.1)
at nuclear density (31014 g cm-3) e.o.s. stiffens and subsonic core collapse slows down
supersonic core continues and “rebounces”: shock wave and SN explosion (“prompt” scenario)
However: unsuccesful ! Shock stalls and eventually recollapses
neutrino losses + iron material dissociation
“delayed” scenario: shock revival by neutrino energy deposition
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Revival of a stalled Supernova shock by neutrino heating Radial trajectories of equal mass shells
- Wilson, Proc. Univ. Illinois, Meeting on Numerical Astrophysics (1982) - Bethe & Wilson, ApJ 295 (1985) 14
Shock formation
Proto Neutron Star
Accretion onto the PNS
Supernova ejecta
Hot bubble
Shock propagation
Neutrino sphereFrom Janka
shock wave
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Rampp & Janka, ApJ 539 (2000) L33
1- D Failed Explosions
Mezzacappa et al., PRL 86 (2001) 1935
Spherically symmetric simulations, Newtonian and General Relativistic, with the most advanced treatment of neutrino transport do not produce explosions.
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prompt e burst
shock breaks through neutrino sphere:
nuclei dissociation
protons liberated allow for quick neutronization
e burst (10-2 s)
Beyond the shock: proto-neutron star (R~30 Km,) which contracts, deleptonizes and cools via all flavor (anti) neutrino emission (10 s)
enpe
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Supernova Neutrinos: Numerical Neutrino Signal
Totani, Sato, Dalhed & Wilson, ApJ 496 (1998) 216
NC CC
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Neutrino flux spectra and overall features
Neutrinos trapped in the high density neutrino-sphere
at the emission surface (R ~ 10-20 Km)
T ~ 2<E>/3 ~ GMmN/3R ~ 10 – 20 MeV
Emission via diffusion
tdiff ~ R2/ ~ GF2 E2 nN ~ 102 cm tdiff =
O(1 s)
Total luminosity
Etot ~ GM2/R ~ 1053 erg
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Neutrino energy distribution
T ~ <E>/3
e <E> ~ 10 –12 MeV
e <E> ~ 14 –17 MeV
, , , <E> ~ 24 –27 MeV
opacity regulated by scattering on (less abundant) protons
opacity regulated by neutral current only
TEe
E
dE
dL/
3
1
1 2 3 4
0.2
0.4
0.6
0.8 Fermi-Dirac-like =2
Maxwell-Boltzmann-like
Equipartition of flux
L(e) ~ L( e) ~ L(x) ~ L( x)
Cross-sections depends on energy; T and density profile
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Time evolution of neutrino signal
prompt e burst 1051 erg in #10 msec
other flavor (anti)neutrino energy and luminosities raises when shock stalls and matter accretes (100 ms) 10% - 25% of the total luminosity in 0.5 sec
Formed protoneutron star cooling 90% -75% of total luminosity
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SN1987A at Kamiokande and IMBSupernova explosion of Sanduleak-69202 in the Large Magellanic Cloud (50 Kpc)
Neutrino observed at Kamiokande II, IMB (water cherenkov) and Baksan (scintillation light) at 7:35:40 UT on 23th february 1987. Optical brightness at 10.38 UT
Detection: KII and IMB
Baksan
ee
eFO
enp
xx
e
e
1616
eNC
enp
e
e
1616
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Time energy analysis
(Loredo and Lamb 1995)
T(t)=Tc0/(1+t/3c)
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SN & ICARUSSN explosion rate
In our galaxy 7.3 h2 per century (from observations in other galaxies)
Large Magellanic Cloud 0.5 per century
but record of hystorical SN suggests a larger number
A rate of 1 per year requires distances of 15 Mpc (Virgo cluster) (too low signal in ICARUS. See later)
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Detection tecnique
- Elastic scattering
Recoil electron direction highly correlated to direction
Larger for e (prompt pulse)
245,
245,
245
245
cm )(103.1)(
cm )(106.1)(
cm )(108.3)(
cm )(102.9)(
MeVEe
MeVEe
MeVEe
MeVEe
e
e
TMeV .6Ktons
1.2Ktons
e e 3.5 4 8e e 5 2 4, e 8 1 2, e 8 1 2total 8 16
ICARUS initial physics program
SN @ d=10Kpc
dtdEdEE
EEdEtL
d
Ndn e
e
ee
),()()(
4 2
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e capture
super allowed Fermi
and GT transitions
*4040 KeAre
rays 40 K
T MeV 0.6ktons
1.2ktons
Fermi
11 15 30
GT 11 30 60
total 45 90Good sensitivity to prompt e
burst and to first 100 ms flux
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caveats: no energy dependent sensitivity and energy threshold
no oscillation effects (some result by Vissani,Cavanna,Palamara Nurzia: full swap)
Similar results in
Thompson et al
2002
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SNO, SK, LVD
SK water Cherenkov detector (32 ktons)
e flux raises after prompt burst
15.4 MeV threshold
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Thompson et al 2002
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SNO D2O detector (1 ktons)
ennD
eppD
pnD
pnD
x
x
xx
xx
Eth 2.2MeV
Eth 1.4 MeV
Eth 4 MeV
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Thompson et al 2002
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LVD scintillator counters
expected events: 102 CC
10 NC
ee
CC
BeC
NeC
nep
xxxx
xxxx
e
e
e
)()(
*)()( 1212
1212
1212
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Oscillations
(under study)
General expectations:
1. Prompt e much harder to observe (reduced x interactions)
2. Harder e flux, due to mixing
3. e , enhances energy transfer from neutrino flux to matter behind the stalled shock
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Issues to be studied• neutrino fluxes as a diagnostic tool for SN model: prompt e burst, 100 ms shock revival and all flavor neutrino fluxes
• ICARUS may be sensible to prompt breakout, O(10) e events, good directionality.
• outlook: neutrino oscillations (trigger design)
detection efficiency
neutrino cross section at 10-80 MeV
SN parameters which may be significantly
distinguished : e.o.s., neutrino oscillations,
density profile, neutrino mass, neutrino-
sphere parameters
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dEdtdEdtttt
ε(EEEdE
dσEP
NtEdd
LtEdN
e
eeA
BA
TBBeA
')'(
) ),( )(
)',( 4
),(2
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Star evolution
Stellar structure
- Hydrostatic equilibrium
- Energy conservation
- Energy transfer
2
)()(
r
rrMG
dr
dp
thermal pressure:
negative specific heat
degeneracy pressure:
positive specific heat
4)( 2r
dr
rdL
gravnucl
dr
aTdrrL
)(
3
4- )(
42
111 e
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