High-Energy Neutrino AstronomyHigh-Energy Neutrino Astronomyand and

Gamma-Ray BurstsGamma-Ray Bursts

Kohta Murase (Kohta Murase ( 村瀬　孔大村瀬　孔大 ))(Yukawa Institute for Theoretical Physics, (Yukawa Institute for Theoretical Physics, 　　 Kyoto Kyoto

University)University)Collaborators: S. Nagataki, K. Ioka, T. Nakamura, F. Iocco, S.D. Serpico, T. Koi, H. Takami, K. Sato, K. Asano, S. Inoue

OutlineOutline

Future Prospects for High-Energy GRB Neutrinos

1. Introduction

2. Prompt Neutrino Emission

3. Other Predictions

4. Implications of the GRB-UHECR Hypothesis

5. Summary

マスタ サブタイトルの書式設定

3

Introduction

Why Neutrinos?Why Neutrinos?

• Astrophysical MeV Neutrinos

Solar neutrinos / Supernovae (SN1987A)

Probe of stellar physics (core / core collapse)

• Astrophysical >0.1 TeV Neutrinos

SNRs, GRBs, AGN, cluster of galaxies…

Probe of CR acceleration (pp/pγ)

CRs → ○ direct probe of CRs ×impossible except for UHECRsγs → ○ easily detected ×contamination by leptonic components νs → ○ good proof of CRs ×difficulty of detection

Cosmic-Ray SpectrumCosmic-Ray Spectrum

Knee ~ 1015.5 eV

2nd Knee ~1017.5 eV

Ankle ~ 1018.5 eV

Problems!•What is the source?•What is origin of breaks?•Composition?•Anisotropy?• • •

•How are CRs accelerated?

What Is the source of CRs?What Is the source of CRs?

)( )101~(eBc

pacct －　

coolacc tt •Acceleration vs Loss/Escape should be considered

Hillas condition E < e B L(“necessary” condition)

Source candidates

•<2nd knee SNRs

•2nd Knee ~ AnkleGalactic WindsHypernovae, AGNs, Clusters…just transition?

•> Ankle AGNs, GRBs, Clusters, Magnetars…

for Bohm limit

“Hillas plot”

GRB

AGN jet

clusters

How Are CRs Accelerated?How Are CRs Accelerated? Particle acceleration 　← collisionless shocks

The most popular mechanism

　→ 11stst order Fermi acceleration mechanism order Fermi acceleration mechanism

rrcc=v1/v2=4 → spectral index p = 2=v1/v2=4 → spectral index p = 2 (NR)

relativistic shock acceleration (e.g., GRBs, AGNs)

• GRB case

internal shocks, reverse shock → Γrel ~ a few

forward shock → 　 Γ ~ 100

If diffusive → spectral index p ~ 2 (p~2.2 for Γ→ ∞)

shock

v1 > v2

Shock acceleration

Physics of particle acceleration is not well understood → need more studies

•amount of accelerated particles?•Amplification of magnetic field over large scales?•Diffusive?, or large-angle scattering rather than small-angle scattering? •Other acceleration mechanisms such as 2nd order Fermi acceleration?

ν detection → implications for the source physicsamount of CRs, implications for B, etc.

baryonic (e.g., fireball) vs non-baryonic (e.g., magnetic) ?

Large Neutrino TelescopesLarge Neutrino Telescopes

taken from IceCube homepage

IceCube IceCube (Antarctica)(Antarctica)

Km3 ice Cherenkov ANTARES

taken from ANTARES homepage

KM3Net KM3Net (Mediterranean(Mediterranean ））Km3 water Cherenkov

Complementary sky coverage!Complementary sky coverage!

•Extension of IceCube (Deep Core for < TeV, IceCube II for VHE νs) Extension of IceCube (Deep Core for < TeV, IceCube II for VHE νs)

+NEMO, NESTOR

Detection of NeutrinosDetection of Neutrinos

νμ + N μ + X

μ emits Cherenkov radiation;

direction reconstructed from correlations between PMTs

106 muons from cosmic rays/muon

from neutrinos

Select only muons from

below

Select only muons from

below

•ν-induced　cascade•τ　double-bang　　　　　　　　　　　　　　　　　　　　　　　　　　　etc...　

Other　signals

Water or Ice Cherenkov detectors

•e.g., Radio detection (RICE, ANITA, ARIANNA)e.g., Radio detection (RICE, ANITA, ARIANNA)•Detection of air showers from earth skimming νDetection of air showers from earth skimming νττs (PAO, TA, Ashra)s (PAO, TA, Ashra)

Afterglows Ep,max ~ ZeV

EeV ν, GeV-TeV γ

Prompt EmissionEp,max ~ EeV-ZeV

PeV ν, GeV-TeV γ

Below Photosphere

Ep,max ~ PeV-EeV

TeV-PeV ν, invisible γ

Meszaros (2001)

Flares/Early AfterglowsEp,max ~ EeV

PeV-EeV ν, GeV γ

マスタ サブタイトルの書式設定

11

Prompt Neutrino Emission

Prompt Emission from Classical (High-Luminosity) GRBs

Internal Shock ModelPeV ν, GeV-TeV γ

(Waxman & Bahcall 97)

Meszaros (2001)

Classical AfterglowsExternal Shock Model

EeV ν, GeV-TeV γ (Waxman & Bahcall 00)

(Dai & Lu 01)(Dermer 02)

(Li, Dai & Lu 02)

Prompt Neutrino EmissionPrompt Neutrino Emission

2.0~κπ+n→Δ→γ+p p+

εp

Cosmic-ray Spectrum (Fermi)

Key parameterCR loading

1018.5eV1020.5eV

εγ

Photon Spectrum (Prompt)

εγ,pk~300keV εmax

Photomeson production efficiency~ effective optical depth for pγ process

fpγ ~ 0.2 nγσpγ (r/Γ)

Δ-resonance Δ-resonance approximation

εp εγ ~ 0.3 Γ2 GeV2

εpb~ 0.3 Γ2/εγ,pk ~ 50 PeV

εp2N(εp)

2-α~1.0

2-β~-0.22-p~0

~ΓGeV

εγ2N(εγ)

EHECR≡εp2N(εp)

~εγ,pk2N(εγ,pk)

multi-pion production

Photomeson Production

)7.04.0(~κX+πN→γ+p p± －

(in proton rest frame)

Δ-resonance approximationpion energy επ~ 0.2 εp

break energy επb~ 0.06 Γ2/εγ,pk ~ 10 PeV

επ

Meson Spectrum

επｂ επ

syn

β-1~1.2

α-1~0

επ2N(επ)

Neutrino Spectrum

ενb

β-1~1.2

α-1~0

εν2N(εν) )(→

)()(e→ ee

meson cooling before decay(meson cooling time) ~ (meson life time)→ break energy in neutrino spectra

Neutrino oscillation(Kashti & Waxman 05)

~fpγEHECR

α-3~-2.0

meson & muon decay

“Waxman-Bahcall” type spectrum (Waxman & Bahcall 97, 99)

ενμsyn εν

πsyn

εν

α-3~-2.0

Δ-resonance approximationneutrino energy εν ~ 0.25 επ ~ 0.05 εp 　

•ν lower break energy ενb ~ 2.5 PeV

•ν higher break energy ενπsyn ~ 25 PeV

0:2:1::e

μμee±

μμ±± ν+ν+)ν(ν+e→)ν(ν+μ→π

1:1:1::e

8.1:8.1:1::e

low εν

high εν

Numerical CalculationNumerical Calculation

•CR cooling → 　 synchrotron, Inverse Compton, adiabatic, Bethe-Heitler, photomeson, pp reaction, photodisintegration (nuclei)

•Treatment of Meson Production photomeson production (experimental data + Geant4) pp reaction (Geant4 + SIBYLL-based formulae)

★Multi-pion production can be important (KM & Nagataki 06)

(flux-enhancement by ~(2-3) for α~1, ~10 for α~0.5)

•Meson cooling → synchrotron, IC, adiabatic, πγ, πp

★Spectrum can be complicated → influence on estimate of events

KM, PRD, 76, 123001　(2007), KM et al. (2008)

Prompt CR AccelerationPrompt CR Acceleration

r~1013-1015.5 cm

•Inner range (~1012-13 cm) pγ efficient, UHECR impossible•Middle range (~1013-14 cm) pγ moderately efficient, UHE proton possible•Outer range (~1015-16 cm) pγ inefficient, UHE nuclei survive

(e.g., KM & Nagataki, 2006)

(r-determination ← GLAST (e.g., KM & Ioka 08, Gupta & Zhang 08, Ioka’s talk)

Fig. fromGuetta (07)

(assumption)

Gyro factor ~ (1-10) → tacc ~ max[tcool, tdyn] ⇔ Emax ~ Z 1019-21 eVWaxman (95)

Prompt Neutrino EmissionPrompt Neutrino Emission

Γ=300, Uγ=UB

Set A: Eγ,iso=1053 ergs, r ~ 1013-14.5 cm → muon events ~ 0.1Set B: Eγ,iso=1053 ergs, r ~ 1014-15.5 cm → muon events ~ 0.01

Set C: Eγ,iso=1054 ergs, r ~ 1013-14.5 cm → muon events ~ 1(Note: C is a very extreme case with α=0.5 and β=1.5)

We expect ν signals from one GRB for only nearby/energetic bursts.

A r~1013.5 cm

B r~1014.5 cm

z=1.0

We will need to see as many GRBs as possible with time- and space-coincidence.

The Cumulative BackgroundThe Cumulative Background

• ~10 events/yr by IceCube ( fiducial baryon load)• The most optimistic model is being constrained by

AMANDA/IceCube group. (Achterberg et al. 08)

fiducial baryon loadingEHECR ~ 0.5 EGRB,γ

(Up=10Uγ)

higher baryon loadingEHECR ~ 2.5 EGRB,γ

(Up=50Uγ)

The key parameter baryon loading 　 ΕHECR ≡εp

2 N(εp)

Set A - r~1013-14.5cm Set B - r~1014-15.5cm

Γ=102.5, Uγ=UB

KM & Nagataki, PRD, 73, 063002　(2006)

Current AMANDA limit

fpγ(EHECR/EGRB,γ)<3 → Towards testing the GRB-UHECR hypothesis via νs!

We cumulate neutrino spectra using GRB rate histories. for GRB rate models

(e.g., Guetta et al. 04, 07)

マスタ サブタイトルの書式設定

19

Other Predictions

Early AfterglowsEeV ν, GeV-TeV γ

(Dermer 07)(KM 07)

Prompt Emissionfrom Low-Luminosity GRBs

PeV ν, GeV-TeV γ(KM et al. 06)

(Gupta & Zhang 07)

Meszaros (2001)

FlaresPeV-EeV ν, GeV γ(KM &Nagataki 06)

Novel Results of Swift (GRB060218)Novel Results of Swift (GRB060218)

1. Low-luminosity (LL) GRBs?• GRB060218 (XRF060218) ・ The 2nd nearby event (~140Mpc)

・ Associated with a SN Ic

・ Thermal component (shock breakout?)

・ Much dimmer than usual GRBs (ELL,γ ~ 1050 ergs ~ 0.001 EGRB,γ)

・ LL GRBs (e.g., XRF060218, GRB980425) more frequent than HL GRBs local Rate ~ 102-3 Gpc-1 yr-3

(Soderberg et al. 06, Liang et al. 07 etc…)

If true → contribution to (UHE)CRs & νs

Liang et al. (07)

dark bright

Luminosity

•Prompt ν emission (KM et al. 06, Gupta & Zhang 07) ~1 event from a LL GRB at 10 Mpc•Interaction with the thermal component (KM et al. (06), Yu, Dai, & Zheng (08)) ~(0.01-1) event from a LL GRB at 10 Mpc

Rate

2. Flares in the early afterglow phase (Chincarini’s talk)

• Energetic (Eflare,γ ~ 0.1 EGRB,γ) (e.g., Falcone et al. 07)

(Eflare,γ ~ EGRB,γ for some flares such as GRB050502b)

• δt >~ 102-3 s, δt/T < 1 → internal dissipation models (e.g. late internal shock model vs magnetic dissipation model)

• Flaring in the (opt/)far-UV/x-ray range (Epk ~ (0.1-1) keV)

• Relatively lower Lorentz factors (maybe) (Γ ~ a few×10)

• Flares are common (at least 1/3-1/2 of LGRBs) (also seen in SGRBs)

・ if baryonic (possibly dirty fireball?) ・ more copious photon field→ neutrinose.g., giant flare at z<~0.1 → ~ a few events

Novel Results of Swift (Flares)Novel Results of Swift (Flares)

GRB 050502bFalcone et al. (05)

EnergeticsEnergetics

Neutrino　Energy　Flux∝ Photomeson (p→π)Production Efficiency

NonthermalBaryon Energy×Rat

e×

HL GRB(Waxman & Bahcall

97)

Flare(Murase & Nagataki

06)

LL GRB(Murase et al. 06)

(Gupta & Zhang 07)

Isotropic energy 1 ~0.01-0.1 0.001

Meson Production

Efficiency

1 10 1

Apparent Rate 1 1 ~100-1000 The contribution to

neutrino background1 ~0.1-1 ~0.1-1

↓Normalizing　all　the　typical　values　for　HL　GRBs　to　1

Hence, we can expect flares and LL GRBs are important!

Neutrino Predictions in the Swift EraNeutrino Predictions in the Swift Era

Possible dominant contribution in the very high energy region

KM & Nagataki, PRL, 97, 051101 (2006)KM, Ioka, Nagataki, & Nakamura, ApJL, 651, L5 (2006)

Approaches to GRBs through high-energy neutrinos LL GRBs→a possible indicator of SNe Ibc associated with LL GRBsFlares→information on flare models (baryonic or nonbaryonic etc.)

νs from LL GRBs → little coincidence with bursts, a few events/yrν flashes → Coincidence with flares/early AGs, a few events/yr

Flares (Eflare,γ = 0.1 EGRB,γ ) LL GRBs

(ELL,γ ~ 0.001 EGRB,γ )

HL GRBs

Gupta & Zhang (2007)

baryon loadingEHECR ~ 0.5 Eγ

KM, PRD, 76, 123001 (2007)

Below Photosphere

TeV ν

(Meszaros & Waxman 01)

(Schneider et al. 02)

( Razzaque, Meszaros, & Waxman 03)

(KM et al. 08)

Meszaros (2001)

Below/Around PhotosphereBelow/Around PhotosphereBelow/around the photosphere (even inside the star)CR acceleration could be expected

(c.f., Meszaros & Waxman 01, Razzaque et al. 03)

Below/around photosphere ⇔ small collision radius r(pp optical depth fpp) ~ 0.5 np σpp(r/Γ) >~ 0.1 → pp reaction important

r=1012.5 cmΓ=100, UB=0.1Uγ

strong meson/muon cooling ↓

kaon-contribution is important (Ando & Beacom 05,Asano & Nagataki 06)

Note: kaon-contribution is just roughly estimated in the left panel.

KM (08) Similar calculations are done by Wang, independently.

Successful & Failed GRBsSuccessful & Failed GRBs

Fig. from Razzaque, Meszaros, & Waxman

termination shock

•internal shocks → dissipation→ particle acceleration & radiation

•termination shock → dissipation→ thermalization ~keV → target photons

•Jet penetration success → GRBs failure → failed GRBs 　　　 (Meszaros & Waxman 01)

•Possible two Contributions (IS+TS)•Meson cooling is important (Razzaque, Meszaros, & Waxman 03)

•Spectrum becomes complicated

•Nearby events → ~(10-100) events

KM et al. (08)

The Cumulative Neutrino BackgroundThe Cumulative Neutrino Background

Iocco et al., ApJ (08), KM et al., in prep. (08)

Schneider et al. 02 (First Stars)

Precursor = successful GRBs (GRB rate)

(c.f. sub-photosphere νs)

•Possible choked-jet signals?Possible choked-jet signals?•choked νs → diffuse background •The AMANDAII limit implies (# of failed GRBs)/(# of SNe-Ibc) < 0.1 for EHECR ~ 0.5 EGRB

•Possible POPIII contribution?Possible POPIII contribution?•Schneider et al. (02) → overestimation•ν detection would be difficult, even if all the first stars can produce GRBs.

マスタ サブタイトルの書式設定

29

Implications of theGRB-UHECR Hypothesis

Test of GRB-UHECR HypothesisTest of GRB-UHECR Hypothesis

3153

UHECRGpcyrergs~ 10E

－－ ergs10E53

,GRB～

possible but requiring high baryon load… EHECR >~ EGRB,γ

(Swift-era lower local rates may lead to higher baryon load…)

•GRBs could be UHECR sources (Waxman 95, Vietri 95)•PAO → possible correlation with galaxies as well as AGN (e.g., Kashti & Waxman 08, Ghisellini et al. 08)

Gpcyr1 31HLGRB

－－～

← Model predictions with UHECR-normalization (except for Late Prompt/Flare)

Comparable

HL Prompt ~ 30 events/yrLL Prompt ~ 10 events/yr Afterglow (ISM) ~ 0.1 events/yr 　Afterglow (WIND) ~ 1 event/yrLate Prompt/Flare ~ 2 events/yr (↑ NOT UHECR sources)

KM, PRD, 76, 123001 (07)

Cosmogenic NeutrinosCosmogenic Neutrinos

Strong evolution model → possible detection in the near future (see Yuksel & Kistler 07)

Takami, KM, Nagataki, and Sato (08)

Auger 07 limit

AssumptionEmax=1022eV

CMB+CIB(Best-fit model of Kneiske et al. 04)

νs generated outside the source by UHECR-CMB/CIB interactions

strong evolution∝(1+z)4.9 (z<1.2)

∝(1+z)0.2 (z>1.2)

normal evolution∝(1+z)3.5 (z<1.2)

∝(1+z)-1.2 (z>1.2)

no evolution∝(1+z)0

Thick: dip modelThin: ankle model

UHECR AstronomyUHECR AstronomyUHECR production may be possible in both of high- and low-luminosity GRBs

The source number density ← PAO, TA Burst Models → 　 Sensitive to effective EGMF strength

(structured + intergalactic)Necessity of future observations and theoretical studies PAO results ~10-4 Mpc-3 → HL GRB marginally inconsistent

(Takami & KM 08)

(KM, Ioka, Nagataki, & Nakamura, PRD, 08)

HE CR acceleration → inevitable γ-ray emission (>TeV)

Internal attenuation due to pair-creationpγ ⇔ γγ 　 (Waxman & Bahcall, PRL, 97, Dermer et al., ApJL, 07)

“Roughly speaking”… νbright (dark) ⇔ TeV γ dark (bright)

Nγ ⇔ γγ (KM, Ioka, Nagataki, & Nakamura, PRD, 08)

Survival of UHE heavy nuclei (e.g., Fe) → τγγ(TeV) <~ 1

•PAO → (tentative) existence of heavier nuclei (Unger et al. 07)

UHE nuclei → large r or Γ (even subdominant) TeV γ rays

•small r or Γ ⇔ ν bright, while TeV γs attenuated→ electromagnetic cascades & inevitable GeV γ rayslarge baryon loads → spectrum modification ⇔ GLAST (c.f., Dermer & Atoyan 04, Asano & Inoue 07)

CR-Induced Gamma-RaysCR-Induced Gamma-Rays

GRB-UHECR hypothesis could be tested in the future…

(see also, Wang et al., ApJ, 08)

SummarySummary• We can expect high-energy neutrino signals under the

internal/external shock models if jets are baryonicif jets are baryonic.

　 1. Prompt ν emission models (HL GRBs or LL GRBs) have been tested Prompt ν emission models (HL GRBs or LL GRBs) have been tested by AMANDA/IceCube.by AMANDA/IceCube.

2. Neutrino flashes (flares) and Neutrino flashes (flares) and neutrinoneutrino early afterglows early afterglows

3. Neutrinos from sub-photospheresNeutrinos from sub-photospheres

• Time- and space-coincidence (with Swift, GLAST etc.)Time- and space-coincidence (with Swift, GLAST etc.) → → more merits than other sources (clusters, AGN etc)more merits than other sources (clusters, AGN etc)

• A good probe of CR acceleration

• Non-detection → just constraints on models

• Detection → clues to GRB models and CR acceleration (poor statistics → importance of multi-messenger astronomy)

• The connection between GRBs and UHECRs will be tested

Waiting to detect extragalacticνs hopefully…Waiting to detect extragalacticνs hopefully… Thank you!

SparesSpares

Photomeson ProductionPhotomeson Production

Δ-resonance

multi-pion production

hadrons→γ+p

•Thresholdεp εγ ~ 0.2 GeV2

(Eγ~ 0.145 GeV)

•Δ-resonanceΔ-resonanceεεpp ε εγγ ~ 0.3 GeV ~ 0.3 GeV22

(Eγ~ 0.32 GeV)

εεππ~ 0.2 ε~ 0.2 εpp

εενν~ 0.25 ε~ 0.25 εππ ~ 0.05 ε ~ 0.05 εpp

•multi-pion production　　　 Eγ >> GeV 　 <Nεπ> ~ 0.5 εp

e.g., GZK mechanism εγ ~ 10-3 eV εp ~ 1020 eV εν~ 5×1018 eV

hadrons→pp

photodisintegration

nor p 1)-(A →A γGiant Dipole Resonance (GDR)

Important for survival of UHE heavy nuclei

Important for dense targets

Ep>>TeVHigh inelasticityHigh multiplicity

Other EffectsOther Effects•Effect of multi-pion production

2.2~βby a factor

by one order

For flatter photon spectra 　(α=0.5,β=1.5)

For typical photon 　 spectra (α=1, β=2.2)

KM & Nagataki, PRD, 73, 063002 (2006)

•Effect of Kaon originated neutrinos

Asano & Nagataki (06)

•Effect of neutrino oscillation

νeνμντ=1:1:1

νeνμντ=1:1.8:1.8

e.g., Kashti & Waxman (05)

EeV neutrinos from Optical/IR FlashesEeV neutrinos from Optical/IR Flashes

Neutrinos can be detected only if we observe very strong optical/IR flashes (when the deceleration radius is small)! Lack of optical/IR flashes (and early-AG may not behave as expected)•Dust extinction? (Roming et al. 05). •Internal dissipation origin?•RS emission (thin ejecta) is fainter than earlier estimates (Nakar & Piran 04)•Suppression due to IC cooling (thick ejecta). (Beloborodov 05)•Highly magnetized flow (Zhang & Kobayashi 05, Luitikov 05)

990123 at z=0.1

z=0.1

KM, PRD, 76, 123001 (2007)

Neutrinos from LL GRBsNeutrinos from LL GRBsEx.) XRF060218-like burst•Prompt nonthrmal emission Epk ~ 5 keV

↑Internal shock model (e.g., Toma, Ioka, Sakamoto, & Nakamura 07)

•Prolonged thermal emission kT ~ 0.15 keV

KM, Ioka, Nagataki, & Nakamura, ApJL, 651, L5 (2006)

r/Γ2=fixed

D=10Mpc

Muon events ~ 1 event

Muon events ~ 0.1 event

See also Gupta & Zhang 07For early afterglowssee Yu, Dai, & Zheng (08)

Early Afterglow Neutrino EmissionEarly Afterglow Neutrino Emissionshallow decay

・ Forward Shock Model (energy injection etc.)・ Reverse Shock Model (Genet et al. 07, Beloborodov 07)

・ Late Prompt Emission Model (Ghisellini et al. 07)

Typically Eshallow,γ ~ 0.1 EGRB,γ (Liang et al. 07)• RS-FS model: ν-detection by IceCube would be difficult…• Late prompt emission model: ν-detection may be possible.

t

Its origin is still controversial...

Prompt

KM, PRD, 76, 123001 (2007)

high baryon loadingEHECR ~ EGRB,γ

fiducial baryon loadingEHECR ~ Eshallow,γ

~ 0.1 EGRB,γ

steep decay

(For forward shock νs,see Dermer 02, 07)

High-Energy Early Afterglow EmissionHigh-Energy Early Afterglow Emission

Plateau Emission ~Late Internal activity?~Plateau Emission ~Late Internal activity?~

Survival of Heavy Nuclei?Survival of Heavy Nuclei?Wang, Razzaque, & Meszaros (2008)

KM, Ioka, Nagataki, & Nakamura, PRD, accepted (2008)

Thick: UHE Fe cannot surviveThin: UHE Fe can survive (p 75% & Fe 25%)

Recent PAO results → (tentative) existence of heavier nuclei (Unger et al. 07)

UHE nuclei production in GRBs (IS, RS, and FS models) and hypernovae → Possible at enough large radii r and/or Γ

Survival of UHE heavy nuclei → neutrino “dark” 　 (τγγ<~ 1) → TeV gamma-ray “bright”

ν γCases UHE Fe can survive (p 75% and Fe 25%)

Slow Jet SNe (Slow Γ GRBs)Slow Jet SNe (Slow Γ GRBs)

Razzaque et al. 04,05Ando & Beacom 05

SNe at ~2-3MpcSNe → 100 events!pp neutrinos, and nus from kaons important

High-Energy Emission MechanismsHigh-Energy Emission Mechanisms

• Leptonic ModelsLeptonic Models　 1. Electron synchrotron 1. Electron synchrotron

2. Synchrotron Self-Compton2. Synchrotron Self-Compton

• Hadronic ModelsHadronic Models3. Proton synchrotron3. Proton synchrotron

4. Neutral pion decay produced 4. Neutral pion decay produced by by photo-meson productionphoto-meson production

5. The contribution from 5. The contribution from electrons+positrons produced electrons+positrons produced by by photo-pair productionphoto-pair production

F∝1/2-p

∝1-p/2

∝(3-p)/2

SSCm SSC

KN maxSSC

m

Synchrotron

SSC

m

e-synch.F

max, pmax,ep-synch.

3 2p

Vietri (97),Totani (98)

Waxman & Bahcall(97), Vietri(98),Bottcher & Dermer (98), Dermer & Atoyan (04)Peer & Waxman (05), Asano & Inoue (07).

e.g. Sari, Piran & Narayan (98)

e.g., Sari & Esin (01), Zhang & Meszaros (01), Guetta & Granot(03), Peer & Waxman (04)

)(/

)(/~

ee

eeL

LY

syn

SSC

BB

BB

(Asano & Inoue 07)

High-Energy Spectra in the Internal Shock ModelHigh-Energy Spectra in the Internal Shock Model

B が強いとき proton signatureが見える。

Up=Uγ のとき ( 控えめ )Proton signature は見えない