Probing neutrino physics with a galactic supernovaamol/talks/technical/2009/dighe-cutapp.pdf · Generation of a shock wave Probing neutrino physics with a galactic supernova – p.5/44

Probing neutrino physics with agalactic supernova

Amol Dighe

Tata Institute of Fundamental Research

Mumbai

CuTAPP 2005, Ringberg Castle, Germany, April 25-29, 2005

Probing neutrino physics with a galactic supernova – p.1/44

Why study a rare event(¿¿ Galactic SN: once in a few decades ??)

For neutrino mixing:Identifying normal / inverted mass hierarchyProbing extremely low values of �13

For SN astrophysics: Talk by R. Tomàs

Pointing to the SN in advanceTracking the shock wave propagation inside SN

For being prepared to observe relevant signals:Water Cherenkov / Ice CherenkovCarbon-based ScintillatorLiquid Ar

For long-term future:A guide for design parameters of future long baselineexperiments [superbeams / neutrino factories]


A Type II supernova

“Onion-shell” structure:


Trapped neutrinos before the collapseNeutrinos trapped inside “neutrinospheres” around � � 1010g/cc.

For �e; ��e

For all �; ��

G. Raffelt


The core collapseGravitational core collapse


The core collapseGravitational core collapse

Generation of a shock wave


Neutrino emission after the collapseNeutronization burst:Shock wave breaks up the nuclei ) e� capture enhanced�e emitted at the �e neutrinosphere.Duration: The first � 10 ms

�e; ��e; ��; ��; �� ; ��



Cooling through neutrino emission: �e; ��e; ��; ��; �� ; ��Duration: About 10 secEmission of 99% of the SN energy in neutrinos

Can be used for “pointing” to the SNin advance. (“Early warning”)





¿¿¿ Explosion ???





¿¿¿ Explosion ???

Talk by R. Tomàs


Initial neutrino spectraNeutrino fluxes:F 0�i = �0E0 (1 + �)1+��(1 + �) � EE0�� exp ��(�+ 1) EE0 �E0, �: in general time dependentKnown properties of the spectra:

Energy hierarchy: E0(�e) < E0(��e) < E0(�x)Spectral pinching: ��i > 2

E0(�e) � 10–12 MeVE0(��e) � 13–16 MeVE0(�x) � 15–25 MeV��i � 2–4

G. G. Raffelt, M. T. Keil, R. Buras, H. T. Janka

and M. Rampp, astro-ph/0303226 10 20 30 40

0.01

0.02

0.03

0.04

0.05

0.06

0.07

E(MeV)


Flavor-dependent neutrino fluxes

10

15

20

25

0 250 500 750

0 1 2 3 4 5 6

0 250 500 750

Time [ms]

10

15

20

25

0 1 2 3 4

⟨E⟩

0 1 2 3 4 5 6

0 1 2 3 4

L [

1052

erg

s-1

]

Time [s]

ν−eν−x

solid line: ��e

dotted line: ��x

Model hE0(�e)i hE0(��e)i hE0(�x)i �0(�e)�0(�x) �0(��e)�0(�x)

Garching (G) 12 15 18 0.8 0.8

Livermore (L) 12 15 24 2.0 1.6

G. G. Raffelt, M. T. Keil, R. Buras, H. T. Janka and M. Rampp, astro-ph/0303226

T. Totani, K. Sato, H. E. Dalhed and J. R. Wilson, Astrophys. J. 496, 216 (1998)


Propagation through matter

core

envelope

ρ=1010

0.1

10 14

12g/cc

ν

SUPERNOVA

VACUUMEARTH

10 km

10 R kpcsun 10000 km

Matter effects on neutrino mixing crucialFlavor conversions at resonances / level crossings


Level crossings during propagationNormal mass hierarchy Inverted mass hierarchy

H resonance: (�m2atm, �13), � � 103 g/ccIn � channel for normal hierarchy, �� channel for inverted hierarchyL resonance: (�m2�, ��), � � 10 g/ccAlways in � channel�m2 hierarchy ) Independent dynamics at resonances


Conversion probability at resonance

1−P

Pf

Pf

1−Pf

Core

Vacuum

Envelope

f

Pf � exp��2 � ; � �m22E sin2 2� os 2� � 1ne dnedr ��1 � 1) Pf � 1) Adiabatic resonance

Landau’1932, Zener’1932L resonance always adiabaticH resonance adiabatic for jUe3j2 >� 10�3,non-adiabatic for jUe3j2 <� 10�5

AD, A. Smirnov, PRD 62, 033007 (2000)Probing neutrino physics with a galactic supernova – p.12/44

Fluxes arriving at the EarthMixture of initial fluxes:F�e = pF 0�e + (1� p)F 0�x ;F��e = �pF 0��e + (1� �p)F 0�x ;4F�x = (1� p)F 0�e + (1� �p)F 0��e + (2 + p+ �p)F 0�x :Survival probabilities in different scenarios:

Case Hierarchy sin2�13 p �p

A Normal Large 0 os2��

B Inverted Large sin2�� 0C Any Small sin2�� os2��

“Small”: sin2�13 <� 10�5, “Large”: sin2�13 >� 10�3.Sensitivity to sin2�13 an order of magnitude better than currentreactor experiments !


SN87A

(Hubble image)

Confirmed the SN coolingmechanism through neutrinos

Number of events too small tosay anything concrete aboutneutrino mixing

Some constraints onSN parameters obtained

J. Arafune, J. Bahcall, V. Barger, M. Fukugita, B. Jegerlehner, M Kachelrieß,

C. Lunardini, D. Marfatia, H. Minakata, F. Neubig, D. Nötzold, H. Nunokawa,

G. G. Raffelt, K. Shiraishi, A. Smirnov, D. Spergel, A. Strumia, H. Suzuki,

R. Tomàs, J. V. F. Valle, B. Wood, T. Yanagida, M. Yoshimura, et al


Detecting a galactic SNEvents expected at Super-Kamiokande with a SN at 10 kpc:��ep! ne+: � 7000 – 12000�e� ! �e�: � 200 – 300�e +16 O ! X + e�: � 150–800

Some useful reactions at other detectors:

Carbon-based scintillator: � + 12C ! � +X + (15.11 MeV)

Liquid Ar: �e + 40Ar ! 40K� + e�


Identifying

neutrino mixing scenario

A ?? B ?? C ??


The task at hand

Measure the spectra, determine the mixing scenario.

A. Bandyopadhyay, AD, S. Choubey, G. Dutta, I. Gil-Botella, S. Goswami, D. Indumathi,

M. Kachelrieß, K. Kar, C. Lunardini, H. Minakata, M. V. N. Murthy, H. Nunokawa, G. Raffelt,

G. Rajasekaran, A. Rubbia, K. Sato, A. Smirnov, K. Takahashi, R. Tomàs, J. Valle, et al

��e


The task at hand

Measure the spectra, determine the mixing scenario.

A. Bandyopadhyay, AD, S. Choubey, G. Dutta, I. Gil-Botella, S. Goswami, D. Indumathi,

M. Kachelrieß, K. Kar, C. Lunardini, H. Minakata, M. V. N. Murthy, H. Nunokawa, G. Raffelt,

G. Rajasekaran, A. Rubbia, K. Sato, A. Smirnov, K. Takahashi, R. Tomàs, J. Valle, et al

C. Lunardini, A. Smirnov, JCAP 0306:009 (2003)

Poorly known initial spectra

Only final ��e spectrumcleanly available (till we haveliquid Ar)

Difficult to find a “clean” ob-servable, i.e. one indepen-dent of some assumptionsabout the initial spectra


Some possible observablesDuring neutronization burst: RB � NCC=NNCCase A ) RB=RB0 � jUe3j2 � 0:05

Broadening of spectra:pinched ! antipinched, i.e. � > 2! � < 2Pinching parameter 3hE2i=4hEi2:Information about the spectrum around the peak

Tail fraction: ftail = N(E > Etail)N(E > Eth)Information about the high-energy part of the spectrum

Difficult to find a “clean” observable, i.e. one independent of someassumptions about the initial spectra.

M. Kachelrieß, C. Lunardini, H. Minakata, H. Nunokawa, G. Raffelt,

K. Sato, A. Smirnov, K. Takahashi, R. Tomàs, J. Valle, et al


Exploiting Earth matter effectsNeutrinos Antineutrinos

10 20 30 40 50 60 70

2.5

5

7.5

10

12.5

15

10 20 30 40 50 60 70

2.5

5

7.5

10

12.5

15

E E

(�e, �x, mixed �) (��e, ��x, mixed ��)

�e ��e


Exploiting Earth matter effectsNeutrinos Antineutrinos

10 20 30 40 50 60 70

2.5

5

7.5

10

12.5

15

10 20 30 40 50 60 70

2.5

5

7.5

10

12.5

15

E E

(�e, �x, mixed �) (��e, ��x, mixed ��)

Total number of events (in general) decreasesCompare signals at two detectors

“Earth effect” oscillations are introducedScenarios B, C for �e, scenarios A, C for ��e


Comparing spectra at multiple detectors

C. Lunardini and A. Smirnov, NPB616:307 (2001)


IceCube as a co-detector with SK/HKIceCube primarily meant for neutrinos with energy >� 150 GeV

The number of Cherenkov photons increases beyond statisticalbackground fluctuations during a SN burst

This signal can be determined to a statistical accuracy of � 0:25%for a SN at 10 kpc.

The Earth effects may change the signal by � 0–10%.

The extent of Earth effectschanges by 3–4 % between theaccretion phase (first 0.5 sec)and the cooling phase.

Absolute calibration not essen-tial

Accretion phase

Cooling phase


Useful SN locations for SK-IC comparison

Earth effects detectable inregions (A) and (B)

AD, M. Keil, G. Raffelt, JCAP 0306:005 (2003)


At a single detector(Identifying Earth oscillation frequency)F��e = sin2�12F 0�x + os2�12F 0��e +�F 0 �A� sin2(�m2�Ly)

(F 0��e � F 0�x) sin 2��12 sin(2��12 � 2�12) (12:5=E)

Oscillation frequency: k� � 2�m2�LThe highest frequency in the “inverse energy” dependence of thespectrum

Completely independent of the primary neutrino spectra

Depends only on solar oscillation parameters, Earth density andthe distance travelled through the Earth


Fourier TransformPower spectrum: GN (k) = 1N ��Pevents eiky��2

Energy resolution is a crucial factor

Scintillation detector: a few thousand events

Water Cherenkov detector: a few ten thousand events

AD, M. Keil, G. Raffelt, JCAP 0306:006 (2003)


Passage through the Earth core

FD�e = sin2 �12 F 0�x + os2 �12 F 0�e + �F 0 7Xi=1 �Ai sin2(kiy=2) ;

i kiy �Ai1 �m=2 � 12 sin(2�12 � 4�m) sin(4� � 4�m) O(!)

2 (�m=2 + � ) os2(� � �m) sin(2�12 � 4�m) sin(2� � 2�m) O(!)

3 (�m + � ) sin(2�12 � 2�m) os4(� � �m) sin(2�m) O(!)

4 � � sin2(2� � 2�m) [ os(2�12 � 4�m)�� 12 sin(2�12 � 2�m) sin(2�m)℄ O(!2)

5 �m 12 sin(2�12 � 2�m) sin2(2� � 2�m) sin(2�m) O(!3)

6 (�m=2� � ) �2 sin(2�12 � 4�m) os(� � �m) sin3(� � �m) O(!3)

7 (�m � � ) sin(2�12 � 2�m) sin4(� � �m) sin(2�m) O(!5)


At a scintillation detector

Passage through the Earth core gives rise to extra peaks.

Model independence of peak positions:


At a water Cherenkov detector

High-k suppression:


Peak identificationStatistical fluctuations may "drown" peaks

Use area under the spectrum for peak identification

Quantifying "efficiency" of a detector


Efficiencies of detectors

High-k suppressionaffects the efficiency ofHK for 35Æ < � < 55Æ.(�: nadir angle)

Large number of eventscompensates forpoorer energy resolution

AD, M. Kachelrieß, G. Raffelt, R. Tomàs, JCAP 0401:004 (2004)


Observation of a Fourier peak in ��e

)Eliminate scenario B independently

of SN models !!!


Shock wave

for identifying

neutrino mixing scenario


Shock wave and level crossings

When shock wave passesthrough a resonance region,adiabatic resonances maybecome non-adiabatic for sometimescenario A ! scenario Cscenario B ! scenario C

May cause sharp changes in thefinal spectra even if the primaryspectra are unchanged /smoothly changing

R. C. Schirato, G. M. Fuller,

astro-ph/0205390

G. L. Fogli, E. Lisi, D. Montanino and

A. Mirizzi, PRD 68, 033005 (2003)


��e Survival probability

20 40 60 80 100 120

0.1

0.2

0.3

0.4

0.5

0.6

0.7

E E EE1b 2b 2a 1a

p_

E�!Shifts right as shock propa-gates to lower densities

Correspondence between densities and energies:

�i = mN�m2atm os 2�132p2GFYeEi � 600 g= m3 os 2�13 25 MeVEi 1Ye

For sharp changes in the density profile: �p = sin2(�a � �b) os2 ��


At a megaton water Cherenkov

“Double dip” feature for hEei“Double peak” feature for hE2e i=hEei2

R. Tomás, M. Kachelrieß, G. Raffelt, AD,

H. T. Janka, L. Scheck, JCAP09(2004)015


Single/double dip

“Single/double dip” robust underneutrino flux modelsmonotonically decreasing average energy

“Single/double dip” visible forsin2 2�13 >� 10�5In �e for normal hierarchyIn ��e for inverted hieratchy


Summary [Detector wishlist]

SN � spectra can help identifying the neutrino mixing scenario:normal / inverted mass hierarchysmall / large �13

A positive identification of the Earth effects rules out scenario B(inverted hierarchy � sin2�13 > 10�3) from ��e or scenario A(normal hierarchy � sin2�13 > 10�3) through �e.

comparison between multiple detectors [SK/MWC, IceCube]identifying earth matter oscillations [SK/MWC, LENA, liquid Ar]

Tracking the shock fronts “in neutrinos”recognizes the presence / absence of a reverse shockconfirms the scenario A [liquid Ar] or scenario B [MWC].determines the times the shocks pass through � � 102–104 g/cc

Talk by R. Tomàs


Extra slides


Splitting events into energy bins

Dip-times energy-bin dependent !!!


Tracking the shock fronts

At t � 4:5 sec, (reverse) shock at �40At t � 7:5 sec, (forward) shock at �40

Multiple energy bins ) the times the shock fronts reach differentdensities of � � 102–104 g/cc


Shock wave at SK


Time evolution of observablesPrimary spectra:F 0i (E) = �ihEii ��ii�(�i) � EhEii��i�1 exp��i EhEii�Number of events:

Nobs = N ��x hE��xi2�2��x �(��x + 2)�(��x) + os2 ��(��eg��e2 � ��xg��x2 )�

where gik = hE��x ik�ki �(�i) h��a; E1hEii�i; E2hEii�i�+ ��a; E3hEii�i; E4hEii�i�i

(Time dependence through the time evolution of E1; E2; E3; E4)Energy moments:

Emobs = N "��x hE��xi2+m�2+m��x �(��x + 2 +m)�(��x) + os2 ��(��eg��e2+m � ��xg��x2+m)#


Time dependence of observables

Model hE0(�e)i hE0(��e)i hE0(�x)i �0(�e)�0(�x) �0(��e)�0(�x)

L 12 15 24 2.0 1.6

G1 12 15 18 0.8 0.8

G2 12 15 15 0.5 0.5


Role of neutrinos in explosion

cooling

Neutrino heating

dM/dt

��

��

Neutrino

70 km200 km

Proto−NeutronStar (n,p)

20 km

p

e

n Stalled Shock

Neutrinosphere��

��

Neutrino heating essential, but not enough

No spherically symmetric (1-D) simulations show robust explosions


Ingredients required for explosion

[ms]

R. Buras, H.-T. Janka, M. Rampp,

K. Kifonidis, astro-ph/0303171

Neutrino heating: higher neutrino opacity

Large scale convenction modes

Stiffer equation of state for the core

Rotation of the star

O. E. Bronson Messer, S. Bruenn, C. Cardall, M. Liebendoerfer,

A. Mezzacappa, W. Raphael Hix, F.-K. Thielemann et alProbing neutrino physics with a galactic supernova – p.44/44

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Probing neutrino physics with a galactic supernovaamol/talks/technical/2009/dighe-cutapp.pdf · Generation of a shock wave Probing neutrino physics with a galactic supernova – p.5/44