SUPERNOVA NEUTRINOS AT

ICARUS

G. Mangano

INFN, Napoli

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

H-R DIAGRAM for M3

H burning

turn-off

growing He core

He core burningHe flash

He and H shell burning

white dwarfs

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 !!

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

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

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

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.

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

Supernova Neutrinos: Numerical Neutrino Signal

Totani, Sato, Dalhed & Wilson, ApJ 496 (1998) 216

NC CC

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

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

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

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

Time energy analysis

(Loredo and Lamb 1995)

T(t)=Tc0/(1+t/3c)

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)

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

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

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

SNO, SK, LVD

SK water Cherenkov detector (32 ktons)

e flux raises after prompt burst

15.4 MeV threshold

Thompson et al 2002

SNO D2O detector (1 ktons)

ennD

eppD

pnD

pnD

x

x

xx

xx

Eth 2.2MeV

Eth 1.4 MeV

Eth 4 MeV

Thompson et al 2002

LVD scintillator counters

expected events: 102 CC

10 NC

ee

CC

BeC

NeC

nep

xxxx

xxxx

e

e

e

)()(

*)()( 1212

1212

1212

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

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

dEdtdEdtttt

ε(EEEdE

dσEP

NtEdd

LtEdN

e

eeA

BA

TBBeA

')'(

) ),( )(

)',( 4

),(2

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