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

Diagnosing Models of Gamma-Ray Bursts through Very High-Energy Gamma-Ray Emission. Kohta Murase Tokyo Institute of Technology Center for Cosmology and AstroParticle Physics, OSU. Collaborators: R. Yamazaki, K. Toma, K. Ioka, S. Nagataki.

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

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  1. Diagnosing Models of Gamma-Ray Bursts through Very High-Energy Gamma-Ray Emission Kohta Murase Tokyo Institute of TechnologyCenter for Cosmology and AstroParticle Physics, OSU Collaborators: R. Yamazaki, K. Toma, K. Ioka, S. Nagataki Deciphering the Ancient Universe with Gamma-Ray Bursts, Kyoto

  2. Content • HE emissiondiscussions 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

  3. Early X-Ray Afterglow Emission • Shallow decay emission: difficult to be explained by the simplest standard afterglow model(Talk by Panaitescu) Chincarini+ 05

  4. 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

  5. Late Prompt Emission Model Ghisellini+ 2007 Late prompt: decelerating jet shallow+normal AG break when q~1/G External shock: standard AG model normal decay

  6. Prior Emission Model Yamazaki 2009 Main jet: prompt after T0~103-4s prompt GRB late optical AG Prior jet: g-ray dim precursor shallow+normal x-ray AG

  7. Prior Emission Model (Contd.) • Assumption (AG onset time of prior jet) < (trigger time T0) • Afterglow F(t) ∝ t-a • t=T+T0F(T)=(T+T0)-a→F(T) ~ const. (T<T0)F(T) ~ T-a (T>T0) consistent with Willingale+ 07

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

  9. External Inverse Compton • Those models naturally predict EIC emission prompt or late prompt qsc “Anisotropic” inverse-Compton emission → Contribution from qsc~0 is suppressed In this talk, we focus on leptonic mechanisms

  10. Predicted Spectrum • Klein-Nishina effect is importantgm2 Eb ~ TeV (gm/103)2 (Eb/MeV) >> EKN ~ Ggm me c2 ~ 50 GeV (t/1000s)-3/4 ∝n(3-p)/2 n F prompt or late prompt n EIC ∝n2-a ∝n2-b KN suppression ∝na-q q=p-1 or p ∝n2-a Eb EKN gm2 Eb

  11. Prior Emission Model (MeV Prompt + FS) • electron distribution = standard AG model • seed photon dist. = observed prompt emission predicted without introducing further parameters KM et al. 10 MNRAS 402 L54 Fermi z=0.3 T0=300s Lg,52=3 Ek,52＝3 ee=0,1 eB=0.01 SSC EIC MAGICII

  12. KM et al. 10 MNRAS 402 L54 EIC duration ~ r(t=T0)/G2c ~ 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~DT~1sPrediction: shallow decay is not expected for such bursts

  13. Late Prompt Model (keV Prompt + FS) • Klein-Nishina effect is importantgm2 Eb ~ 0.1 GeV (gm/300)2 (Eb/keV) << EKN ~ Ggm me c2 ~ 10 GeV (t/1000s)-3/4 • SSC from FS will also contribute to HE emission EcSSC ~ gc2 Ec ~ TeV (t/1000s)-1/4 n F late prompt n2-b n AG n2-a n(3-p)/2 SSC n(3-p)/2 n1-p/2 EIC n-(3-p)/2 n1-p/2 na-q q=p-1 or p n2-a Ec Eb EKN gc2 Ec KM et al. 2010b, in prep. Ｆｅｒｍｉ ｒａｎｇｅ

  14. EIC from Two-Component Models • 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-TeVe.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)

  15. Connection to Fermi GRBs? • So far, GeV emission observed by Fermi may be explained by synchrotron emission in the standard ext. shock model Ghisellini+ 10 MNRAS (Kumar & Duran 09, Ghisellini+ 10 Wang+ 10, talk by Meszaros, Piran) • Fermi bursts themselves do not seem to require models for shallow decay emission

  16. Synchrotron and SSC emission？ • Radiative AG (e.g., ee, eB~0.1-1， n~1cc-1) (Ghisellini+ 10) • Adiabatic AG (e.g., eB~10-4, n~10-3 cc-1) (Kumar and Duran 09) • Unless Y >> 1, it is possible to find parameters where Ecut is observedEcut ~ G (h/2p) (6pe2/sTmec)h-1 ~ G 160 MeV h-1 n F n SSC Ｙ Synch. E＊ Ecut EpkSC EKN

  17. Synchrotron Cutoff by IACTs? • Ecut only depends on G except acc. coff. h • In the adiabatic case, Ecut can be seen n F n EKN ee=0.1 eB=10-5 p=2.4 z=1 SSC Synch. E＊ Ecut Ecut E＊ EKN EpkSC KM & Yamazaki 2010 Ecut observation → measurement of evolution ofG

  18. Summary VHE obs.@>10GeV are relevant for diagnosing GRB models • EIC as a diagnosis of multi-component modelsVHE 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, ｆor signals from UHE nuclei)

  19. Synchrotron Cutoff? • Ecut only depends on G except acc. coff. • For appropriate ee/eB, Ecut may be seen n F n ee=0.003 eB=0.001 p=2.1 EKN SSC Synch. EpkSC E＊ Ecut E＊ Ecut EpkSC EKN Ecut observation → measurement of G!

  20. Issue

  21. Emission Mechanisms • Leptonic mechanismsynchrotronsynchrotron self-Componexternal inverse Compton • Hadronic mechanismpgppnuclear de-excitation

  22. Various 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 07) 　　・Existence of slow tail of ejecta leads to a long-lasting RS

  23. 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

  24. Plateau Emission ~Late Internal activity?~

  25. High-Energy Spectra from Afterglows ISMモデル 100s → 10000s → 1000000s WINDモデル 100s → 10000s → 1000000s

  26. Early Afterglows in the Swift era Energy injection Time-dependent parameters dE/dt ∝t^-0.5 εe ∝t^0.4 z=1

  27. High-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

  28. 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)

  29. (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 (EGRBg ~ 1051 ergs ~ 0.01 EGRBg,iso, qjet ~ 0.1 rad) • Related to the deaths of massive stars (association with SNe Ic) variability~ ms Luminosity Afterglow X-ray、optical、radio Prompt (GRB) Gamma-ray～300 keV Duration: a few s～103s Time 10-102s 103-104s

  30. Internal-External Shock Model(Baryonic Jet Model) r ~ 1014 cm r > 1016 cm Interstellar Medium Central Engine Lorent Factor G>100 Luminosity Bulk kinetic energy ↓ Shock dissipation acceleration magnetic fieldheat Time

  31. Prompt Gamma-Ray Emission g-ray emission ⇔ radiation from electrons accelerated at mildly relativistic (Γrel ~ a few) internal shocks Protons may also be accelerated as well as electrons Amati et al. (2002) Isotropic energy Eγiso ~ 1053 ergs broken power-law spectrum N(eg) ∝ eg-a (e<eg,pk) N(eg) ∝ eg-b (e>eg,pk)

  32. Classical Optically Thin Synchrotron Scenario Fig. from Guetta (07) Optically thick ← → Optically thin rph ~ 1012.5 cm rdec ~ 1016 cm • 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

  33. Cosmic-Ray Acceleration in GRBs assumption necessary for UHECRs η~ (1-10) Acceleration time scale Cooling time scale only if tcool ~ tsyn Criterion for acceleration tacc < max[tcool, tdyn] Escape: tdyn < tcool Ep,max = Esyn ~ 1020-21 eV UHECR production is possible For nuclei survival → EO,max = Eog ~ 1016-17 eV r = 1014 cm Waxman (95) E/Γ

  34. Internal-External Shock Model(Baryonic Jet Model) r ~ 1014 cm r > 1016 cm Interstellar Medium Central Engine Lorent Factor G>100 Luminosity Bulk kinetic energy ↓ Shock dissipation acceleration magnetic fieldheat Time

  35. Basics of Prompt Neutrino Emission Cosmic-ray Spectrum (Fermi) Photon Spectrum (Prompt) Key parameter CR loading εγ2N(εγ) εp2N(εp) 2-p~0 2-β~-0 EHECR≡εp2N(εp) ~εγ,pk2N(εγ,pk) 2-α~1.0 total ECR~20EHECR εp εγ ~ΓGeV 1018.5eV 1020.5eV εγ,pk~300keV εmax Photomeson Production Δ-resonance Δ-resonance approximation εp εγ ~ 0.3 Γ2 GeV2 εpb~ 0.3 Γ2/εγ,pk ~ 50 PeV multi-pion production Photomeson production efficiency ~ effective optical depth for pγ process fpγ ~ 0.2 nγσpγ (r/Γ) (in proton rest frame)

  36. Meson Spectrum pion energy επ~ 0.2 εp break energy επb~ 0.06 Γ2/εγ,pk ~ 10 PeV επ2N(επ) α-1~0 ~fpγEHECR β-1~1 For charged mesons → sync. cooling (meson cooling time) ~ (meson life time) → break energy in neutrino spectra α-3~-2.0 επ επｂ επsyn Waxman & Bahcall, PRL (1997) Neutrino Spectrum Gamma-Ray Spectrum εν2N(εν) εg2N(εg) α-1~0 α-1~0 β-1~1 β-1~1 α-3~-2.0 εν εg ενb ενπsyn εgb εgmax • neutrino energy εν~ 0.25 επ ~0.05 εp • ν lower break energy ενb ~ 2.5 PeV • ν higherbreak energy ενπsyn ~ 25 PeV • g-ray energy εg~ 0.5 επ ~0.1 εp • γ lower break energy εgb ~ 5 PeV • γ maximum energy εgmax ~ 0.1 εpmax

  37. Prompt Neutrino Emission KM & Nagataki, PRD, 73, 063002(2006) z=1.0 A r~1013.5 cm B r~1014.5 cm Γ=300, Uγ=UB Set A: EGRBg,iso=1053 ergs, r ~ 1013-14.5 cm → muon events ~ 0.1 Set B: EGRBg,iso=1053 ergs, r ~ 1014-15.5 cm → muon events ~ 0.01 Set C: EGRBg,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. We will need to see as many GRBs as possible with time- and space-coincidence.

  38. The Cumulative Background for GRB rate models (e.g., Guetta et al. 04, 07) We cumulate neutrino spectra using GRB rate histories. KM & Nagataki, PRD, 73, 063002(2006) • ~10 events/yr by IceCube (moderate CR loading) • The most optimistic model is being constrained by AMANDA/IceCube group. (Achterberg et al. 07,08) The key parameter CR loadingΕHECR ≡εp2 N(εp) Γ=102.5, Ug=UB Current AMANDA limit high CR loading EHECR ~ 2.5 EGRBg (Up=50Ug) moderate CR loading EHECR ~ 0.5 EGRBg (Up=10Ug) Set A - r~1013-14.5cm Set B - r~1014-15.5cm fpg(EHECR/EGRBg)<3 → Towards testing the GRB-UHECR hypothesis via νs

  39. Alternative Scenarios? r~1013-1015.5 cm Fig. from Guetta (07) The optically thin synchrotron scenario has several problems e.g., epk-Liso correlation, low-energy index problem… • Alternative scenarios • Photospheric: Emission from the photosphere (t~1, r~1012.5cm) • SSC: Emission from around the deceleration radius (r~1016cm)

  40. The Cumulative Background KM, PRD(R), 78, 101302 (2008) Photospheric: TeV nus from pp (detectable even for h >> 1) • Important probe of dissipation/acceleration below/around rph • The most efficient case (min[fpg,1]~1) SSC: EeV nus from pg (because of optical synchrotron photons) CR loading EHECR ~ EGRBg~ 1051 ergs (for prompt emission) Photospheric ~ 20 events/yr Classical ~ 10 events/yr SSC ~ 0.1 events/yr

  41. Remarks • Key parameters： CR loading EHECR (UHECR hypothesis → EHECR ~ 1-10 EGRBg) Emission radius r (depending on scenarios) • Gamma rays should be but more complicated! pair creation in the source contribution from leptonic components

  42. GeV Gamma Rays Relative small r → VHE g rays (e.g., ｆｒｏｍ p0) cannot escape r ~ 1014 cm p 100% r~1014 cm EHECR/EGRBg = 0.05 r~1014 cm EHECR/EGRBg =5 EHECR/EGRBg = １.5 EHECR/EGRBg = 0.5 Asano & Inoue (2007) Asano, Inoue, & Meszaros (2008) *Here e index (pe=3) is assumed to be steeper than p index (pp=2) EM cascades in the source (modification for high CR loading) GeV g rays → Fermi, MAGIC (e.g., possibly GRB 090510B)

  43. TeV Gamma Rays Relative large r → VHE g rays (e.g., ｆｒｏｍ p0) can escape r ~ 1015 cm (HL GRB) EHECR/EGRBg=1 r ~ 1016 cm (LL GRB) EHECR/EGRBg = 0.5 KM, Ioka, Nagataki, & Nakamura, PRD (2008) *p0g rays are attenuated by CMB (their detection is not easy) Non-cascades in the source (CR synch. emission can be important) TeV g rays → MAGIC, VERITAS (for nearby/energetic GRBs)

  44. Remarks 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)

  45. Early AfterglowsEeV ν, GeV-TeV γ (Dermer 07)(KM 07) Meszaros (2001) Classical AfterglowsExternal Shock ModelEeV ν, GeV-TeV γ (Waxman & Bahcall 00)(Dai & Lu 01)(Dermer 02)(Li, Dai & Lu 02)

  46. Reverse-Forward Shock Model Γ~ 100-1000 afterglow Reverse shock Forward shock ejecta CBM

  47. Forward Shock vs Reverse Shock • Forward-shock acceleration of protons (Dermer 02) Ultra-relativistic shock For typical parameters, Emax ~ Z 1015eV BISM,-6 (t/104 s)-1/8 Very strong amplification of upstream B is required UHECR acceleration at G >> 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 shock The 1st order Fermi acceleration seems possible It might relatively easy to produce UHECRs UHECRs + optical/IR photons (~ T ~ 100 s) → EeV neutrinos

  48. UHECRs and GRBs

  49. Photospheric Emission Scenario e.g., Meszaros & Rees (00), Rees & Meszaros (05) Epeak ~ kT (characterized) Photosphere kT ~ 100keV Peer, Meszaros, & Rees (06) Significance of thermal emission (r<rph) → High radiation efficiency Dissipation/acceleration occurs below/around the photosphere Nonthermal component comes from electrons at r ~ rph

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