Diagnosing Models of Gamma-Ray Diagnosing Models of Gamma-Ray Bursts through Very High-Energy Bursts through Very High-Energy

Gamma-Ray EmissionGamma-Ray Emission

Kohta MuraseKohta Murase Tokyo Institute of TechnologyTokyo Institute of Technology

Center for Cosmology and AstroParticle Physics, OSUCenter for Cosmology and AstroParticle Physics, OSU

Deciphering the Ancient Universe with Gamma-Ray Bursts, KyotoDeciphering the Ancient Universe with Gamma-Ray Bursts, Kyoto

Collaborators: R. Yamazaki, K. Toma, K. Ioka, S. Nagataki

ContentContent• HE emission

discussions motivated by recent Fermi results+ delayed onset, extra component etc.

many models including int.- and ext.- shocks have been discussed leptonic (talks by Meszaros, Dermer, Piran, Wang)hadronic (talks by Meszaros, Dermer, Ioka, Asano)

• Here, I will talk about HE emission at late time from a different motivation

Early X-Ray Afterglow EmissionEarly X-Ray Afterglow Emission

• Shallow decay emission: difficult to be explained by the simplest standard afterglow model(Talk by Panaitescu)

Chincarini+ 05

Many models have been suggested so far…energy injection, time-dependent parameters, long-lasting RS etc.

Multi-component models (e.g., Granot et al. 06, Toma et al. 06, Ghisellini et al. 07, Yamazaki 09)

have been more and more discussed recently

Ex.: two-component model fits by Ghisellini et al. 09

Late Prompt Emission ModelLate Prompt Emission Model

Late promptLate prompt::

decelerating jet

shallow+normal AG

break when ~1/ External shock:External shock:

standard AG model

normal decay

Ghisellini+ 2007

Prior Emission ModelPrior Emission Model

Main jet:Main jet:

prompt after T0~103-4s

prompt GRB

late optical AG

Prior jetPrior jet::

-ray dim precursor

shallow+normal x-ray AG

Yamazaki 2009

Prior Emission Model (Contd.)Prior Emission Model (Contd.)

• Assumption

　 (AG onset time of prior jet) < (trigger time T0)

• Afterglow

F(t) t∝

• t=T+T0

F(T)=(T+T0)

→F(T) ~ const. (T<T0)　 F(T) ~ T (T>T0)

consistent with Willingale+ 07

• Motivated by recent interpretations for x-ray afterglows, let us consider consequences of such two-component models for high-energy emission

External Inverse ComptonExternal Inverse Compton

• Those models naturally predict EIC emission

“Anisotropic” inverse-Compton emission→ Contribution from sc~0 is suppressed

sc

In this talk, we focus on leptonic mechanisms

prompt or late prompt

Predicted SpectrumPredicted Spectrum

• Klein-Nishina effect is importantm

2 Eb ~ TeV (m/103)2 (Eb/MeV)

>> EKN ~ m me c2 ~ 50 GeV (t/1000s)-3/4

∝2-

∝(3-p)/2prompt or late prompt

EIC

F

Eb EKN m2 Eb

∝2-

∝2-∝-q

KN suppression

q=p-1 or p

Prior Emission Model (MeV Prompt + FS)Prior Emission Model (MeV Prompt + FS)

• electron distribution = standard AG model

• seed photon dist. = observed prompt emission predicted without introducing further parameters

z=0.3T0=300sL=3E ＝ 3e=0,1B=0.01

KM et al. 10 MNRAS 402 L54

MAGICII

EIC

SSC

Fermi

EIC duration ~ r(t=T0)/2c ~ T0 ~ 1000 s → Follow-up obs. by IACTs would be possible (~ dozens of seconds)

＊ ~GeV extra comp. of observed Fermi GRBs may be explained for T0~T~1sPrediction: shallow decay is not expected for such bursts

KM et al. 10 MNRAS 402 L54

Late Prompt Model (keV Prompt + FS)Late Prompt Model (keV Prompt + FS)

2-

(3-p)/2

late prompt

EIC

F

Eb EKN c2 Ec

2-

2-

-q

q=p-1 or p

SSC

-(3-p 1-p 1-p/2

Ec

• Klein-Nishina effect is importantm

2 Eb ~ 0.1 GeV (m/300)2 (Eb/keV) << EKN ~ m me c2 ~ 10 GeV (t/1000s)-3/4

• SSC from FS will also contribute to HE emission Ec

SSC ~ c2 Ec ~ TeV (t/1000s)-1/4

(3-p)/2

Ｆｅｒｍｉ ｒａｎｇｅKM et al. 2010b, in prep.

AG

Useful for testing these kinds of two-component models, and quantitative studies of obs. may allow us to discern various theoretical possibilities

Such EIC emission may similarly be expected in such two-component models for prompt emission- MeV prompt + FS/RS (prior emission model)small T0 → extra comp. at GeV-TeV 　e.g., MeV prompt + IS, Toma, Wu & Meszaros 2010

As was previously suggested ， EIC may also lead to GeV-TeV flares or GeV-TeV flashes from RS(e.g., Wang, Li, & Meszaros 2006)

EIC from Two-Component ModelsEIC from Two-Component Models

Connection to Fermi GRBs?Connection to Fermi GRBs?

• So far, GeV emission observed by Fermi may be explained by synchrotron emission in the standard ext. shock model

•Fermi bursts themselves do not seem to require models for shallow decay emission

Ghisellini+ 10 MNRAS(Kumar & Duran 09, Ghisellini+ 10Wang+ 10, talk by Meszaros, Piran)

Synchrotron and SSC emissionSynchrotron and SSC emission ？？• Radiative AG (e.g., e, B~0.1-1 ， n~1cc-1) (Ghisellini+ 10)

• Adiabatic AG (e.g.,B~10-4, n~10-3 cc-1) (Kumar and Duran 09)

• Unless Y >> 1, it is possible to find parameters where Ecu

t is observedEcut ~ (h/2) (6e2/Tmec)-1 ~ 160 MeV -1

Synch.

SSCF

Ecut EKN Epk

SC

E ＊

Ｙ

Synchrotron Cutoff by IACTs?Synchrotron Cutoff by IACTs?

• Ecut only depends on except acc. coff. • In the adiabatic case, Ecut can be seen

Synch. SSC

F

Ecut EKN Epk

SC E

＊

Ecut observation → measurement of evolution of

Ecut

E ＊

EKN

e=0.1B=10-5

p=2.4z=1

KM & Yamazaki 2010

SummarySummaryVHE obs.@>10GeV are relevant for diagnosing GRB models• EIC as a diagnosis of multi-component models

VHE observations at ~102-104 s- prior emission model for shallow decay- late prompt emission for shallow decay etc.

• Syn. cutoff or extra components (SSC or hadronic)VHE observations at ~1-102 s for Fermi GeV bursts - e.g., adiabatic AG or radiative AG models

Maybe difficult by Fermi　 IACTs are better in sensitivities though det. prob. is not large fast follow-up (<100s) & LE thr. (~10GeV) required →CTA (see also my postar #63, ｆ or signals from UHE nuclei)

Synchrotron Cutoff?Synchrotron Cutoff?

• Ecut only depends on except acc. coff.• For appropriate e/B, Ecut may be seen

Synch. SSC

F

Ecut EKN Epk

SC

E ＊

Ecut observation → measurement of

Ecut

E ＊

EpkSC

EKN

e=0.003B=0.001p=2.1

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Issue

Emission MechanismsEmission Mechanisms

• Leptonic mechanismsynchrotronsynchrotron self-Componexternal inverse Compton

• Hadronic mechanismpppnuclear de-excitation

Various InterpretationsVarious Interpretations

Many possibilities have been suggested…

For example,Modified Forward Shock Models a. energy injection (e.g., Sari & Meszaros 00) b. time-dependent parameters (e.g., Ioka et al. 06) 　 c. complicated density profile (e.g., Ioka et al. 06)

Long-lasting Reverse Shock Model(Genet et al. 07, Uhm & Beloborodov 0

7)　　・ Existence of slow tail of ejecta leads to a long-lasting RS

Multi-component models (e.g., Granot et al. 06, Toma et al. 06, Ghisellini et al. 07, Yamazaki

09)

more and more discussed recently

Ex.: two-component model fits by Ghisellini et al. 09

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

High-Energy Spectra from AfterglowsHigh-Energy Spectra from Afterglows

ISM モデル

WIND モデル

100s → 10000s → 1000000s

100s → 10000s → 1000000s

Early Afterglows in the Swift eraEarly Afterglows in the Swift era

Energy injection Time-dependent parametersεe ∝ 　 t^0.4dE/dt ∝ 　 t^-0.5

z=1

High-Energy Gamma-Rays from FlaresHigh-Energy Gamma-Rays from Flares

フレアの high-energy をうけるには近傍のバーストに限られる

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

2. Flares in the early afterglow phase• Energetic (Eflare,γ ~ 0.1 EGRB,γ (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)• Flaring in the (far-UV)/x-ray range (Epeak ~ (0.1-1) keV)• (Maybe) relatively lower Lorentz factors (Γ ~ a few×10)• Flares are common (at least 1/3 ~ 1/2 of LGRBs) (even for SGRBs)

Baryonic (possibly dirty fireball?) vs non-baryonic?↑neutrinos!

Flares

Burrows et al. (07)

(Long) Gamma-Ray Bursts(Long) Gamma-Ray Bursts

•The most violent phenomena in the universe (L~1051-52 ergs s-1)•Cosmological events (z~1-3)•~1000 per year (⇔ apparent rate of ~ 1/10000 of SNe Ibc rate)•Jet hypothesis (EGRB~ 1051 ergs ~ 0.01 EGRB,iso, jet ~ 0.1 rad)•Related to the deaths of massive stars (association with SNe Ic)

Gamma-ray ～ 300 keVDuration: a few s ～ 103s

Prompt (GRB)

AfterglowX-ray 、 optical 、 radio

variability~ ms

Time

Luminosity

10-102s103-104s

Internal-External Shock Model(Baryonic Jet Model)

Lorent Factor

>100

InterstellarMedium

Bulk kinetic energy↓

Shock dissipation　　　　　　

acceleration magnetic field 　heat 　　　

CentralEngine

Time

Luminosity

r ~ 1014 cm r > 1016 cm

Prompt Gamma-Ray EmissionPrompt Gamma-Ray Emission

-ray emission ⇔ radiation from electrons accelerated at

mildly relativistic (Γrel ~ a few) internal shocks Protons may also be accelerated as well as electrons

1~α　 2.2~

Amati et al. (2002)

keV300~pk,

broken power-law spectrum

N(∝,pk

N(∝,pk

Isotropic energyEγ

iso ~ 1053 ergs

•Peak energy of ~ 300 keV is identified with synchrotron peak•The typical required magnetic field is B ~ 104-5 G for Γ ~ 300•The typical emission radius is r~1013-1015.5 cm

Fig. fromGuetta (07)

Classical Optically Thin Synchrotron ScenarioClassical Optically Thin Synchrotron Scenario

Optically thick ← → Optically thin

rph ~ 1012.5 cm rdec ~ 1016 cm

Cosmic-Ray Acceleration in GRBsCosmic-Ray Acceleration in GRBs assumption

necessary for UHECRs

η~ (1-10)

Criterion for acceleration

tacc < max[tcool, tdyn] Escape: tdyn < tcool

Ep,max = Esyn ~ 1020-21 eVUHECR production is possible

For nuclei survival→ EO,max = Eo ~ 1016-17 eV

E/Γ

Waxman (95)

Acceleration time scale

Cooling time scale only if tcool ~ tsyn

r = 1014 cm

Internal-External Shock Model(Baryonic Jet Model)

Lorent Factor

>100

InterstellarMedium

Bulk kinetic energy↓

Shock dissipation　　　　　　

acceleration magnetic field 　heat 　　　

CentralEngine

Time

Luminosity

r ~ 1014 cm r > 1016 cm

Basics of Prompt Neutrino EmissionBasics of Prompt 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~-02-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)

total ECR~20EHECR

pion energy επ~ 0.2 εp

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

επ

Meson Spectrum

επｂ επ

syn

β-1~1

α-1~0

επ2N(επ)

Neutrino Spectrum

ενb

β-1~1

α-1~0

εν2N(εν)

)(→

)()(e→ ee

For charged mesons → sync. cooling(meson cooling time) ~ (meson life time)→ break energy in neutrino spectra

~fpγEHECR

α-3~-2.0

ενπsyn

εν

α-3~-2.0

neutrino energy εν ~ 0.25 επ ~ 0.05 εp 　

•ν lower break energy ενb ~ 2.5 PeV

•ν higher break energy ενπsyn ~ 25 PeV

→0

Gamma-Ray Spectrum

εb

β-1~1

α-1~0

ε2N(ε)

εmax

ε

-ray energy ε ~ 0.5 επ ~ 0.1 εp 　

•γ lower break energy εb ~ 5 PeV

•γ maximum energy εmax ~ 0.1 εp

max

Waxman & Bahcall, PRL (1997)

Prompt Neutrino EmissionPrompt Neutrino Emission

Γ=300, Uγ=UB

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

Set C: EGRB,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.

KM & Nagataki, PRD, 73, 063002　(2006)

The Cumulative BackgroundThe Cumulative Background

• ~10 events/yr by IceCube (moderate CR loading)• The most optimistic model is being constrained by

AMANDA/IceCube group. (Achterberg et al. 07,08)

moderate CR loadingEHECR ~ 0.5 EGRB

(Up=10U)

high CR loadingEHECR ~ 2.5 EGRB

(Up=50U)

The key parameter CR loading 　 ΕHECR ≡εp

2 N(εp)

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

Γ=102.5, U=UB

KM & Nagataki, PRD, 73, 063002　(2006)

Current AMANDA limit

fp(EHECR/EGRB)<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)

r~1013-1015.5 cm

Alternative scenarios •Photospheric: Emission from the photosphere (rcm)•SSC: Emission from around the deceleration radius (r~1016cm)

Fig. fromGuetta (07)

Alternative Scenarios?Alternative Scenarios?

The optically thin synchrotron scenario has several problems e.g., pk-Liso correlation, low-energy index problem…

The Cumulative BackgroundThe Cumulative Background

Photospheric: TeV nus from pp (detectable even for >> 1)• Important probe of dissipation/acceleration below/around rph

• The most efficient case (min[fp,1]~1)SSC: EeV nus from p (because of optical synchrotron photons)

CR loadingEHECR ~ EGRB~ 1051 ergs(for prompt emission)

Photospheric~ 20 events/yrClassical~ 10 events/yrSSC~ 0.1 events/yr

KM, PRD(R), 78, 101302 (2008)

RemarksRemarks

• Key parameters ： 　 CR loading EHECR

(UHECR hypothesis → EHECR ~ 1-10 EGRB)

Emission radius r

　 (depending on scenarios)

• Gamma rays should be but more complicated!

pair creation in the source

contribution from leptonic components

GeV Gamma RaysGeV Gamma Rays

Relative small r → VHE rays (e.g., ｆｒｏｍ 0) cannot escape

r ~ 1014 cmp 100%

Asano, Inoue, & Meszaros (2008)

*Here e index (pe=3) is assumed to be steeper than p index (pp=2)

r~1014 cmEHECR/EGRB = 0.05

r~1014 cm

Asano & Inoue (2007)

EHECR/EGRB =5

EHECR/EGRB = 0.5

EHECR/EGRB = １ .5

EM cascades in the source (modification for high CR loading)

GeV rays → Fermi, MAGIC (e.g., possibly GRB 090510B)

TeV Gamma RaysTeV Gamma Rays

Non-cascades in the source (CR synch. emission can be important)TeV rays → MAGIC, VERITAS (for nearby/energetic GRBs)

r ~ 1015 cm (HL GRB)EHECR/EGRB =1

r ~ 1016 cm (LL GRB)EHECR/EGRB = 0.5

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

Relative large r → VHE rays (e.g., ｆｒｏｍ 0) can escape

*0 rays are attenuated by CMB (their detection is not easy)

RemarksRemarks

CR acceleration during the prompt phase is testable

But prompt emission mechanism is highly uncertain(magnetic dissipation models → less neutrinos…)

Even if prompt emission is magnetic, GRBs can still be candidates of the UHECR origin

(But large Ekin w. small fe is required)

Because CRs are likely to be accelerated in afterglows caused by shock dissipation

(This situation is similar to AGNs)

Early AfterglowsEeV ν, GeV-TeV γ

(Dermer 07)(KM 07)

Meszaros (2001)

Classical AfterglowsExternal Shock Model

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

(Dai & Lu 01)(Dermer 02)

(Li, Dai & Lu 02)

Reverse-Forward Shock ModelReverse-Forward Shock Model

afterglow

Reverse shock

Forward shock

ejecta

Γ~ 100-1000

CBM

Forward Shock vs Reverse ShockForward Shock vs Reverse Shock

• Forward-shock acceleration of protons (Dermer 02)

Ultra-relativistic shockFor typical parameters, Emax ~ Z 1015eV BISM,-6 (t/104 s)-1/8

Very strong amplification of upstream B is requiredUHECR acceleration at >> 1 shock is theoretically difficult→Other mechanisms such as the 2nd order Fermi acceleration?

• Reverse-shock acceleration of protons (Waxman & Bahcall 00)

Mildly relativistic or non-relativistic shockThe 1st order Fermi acceleration seems possibleIt might relatively easy to produce UHECRsUHECRs + optical/IR photons (~ T ~ 100 s) → EeV neutrinos

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UHECRs and GRBs

Photospheric Emission ScenarioPhotospheric Emission Scenario

Significance of thermal emission (r<rph) → High radiation efficiencyDissipation/acceleration occurs below/around the photosphereNonthermal component comes from electrons at r ~ rph

kT ~100keV

e.g., Meszaros & Rees (00), Rees & Meszaros (05)Epeak ~ kT (characterized)

Photosphere

Peer, Meszaros, & Rees (06)

Synchrotron Self-Compton ScenarioSynchrotron Self-Compton Scenario

• Some observations of early AGs lead to r ~ 1016 cm

difficulty in synchrotron scenario (Kumar & MacMahon 08)

• e.g., one-zone interpretation of 080319B → SSC model

optical emission implies large r (e.g., Racusin et al., Nature 08)

opt

~ 300 keV

~ 10 GeVνFν

ν

bright optical!

Racusin et al. (08)

synchrotron

B ~ 100 G

(Asano & Inoue 07)

High-Energy Spectra in the Internal Shock Model

UB>>U=Up

proton signature

Up=U=UB

No proton signature

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

r~1013-1015.5 cm

•Inner range (~1011-13 cm) pγ efficient, UHECR impossible•Middle range (~1013-15 cm) pγ moderately efficient, UHE proton possible•Outer range (~1015-16 cm) pγ inefficient, UHE nuclei survive

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

r-determination is important ← GLAST (e.g., KM & Ioka 08, Gupta & Zhang 08)

Fig. fromGuetta (07)

Issues of Prompt EmissionIssues of Prompt Emission

DO not belive the synchrotron scenario.

Magnetic Dissipation - External Shock Model(Magnetic Jet Model)

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)