meszaros

Slide1: 

Peter Mészáros Pennsylvania State University Astrophysical Sources of TeV to GZK neutrinos

Slide2: 

Energy (eV) Radio CMB Visible GeV g-rays 1 TeV Flux 400 microwave photons per cm3 [slides: Halzen 03]

Slide3: 

TeV sources! cosmic rays / / / / / / / / / / / / / / / / / n

Slide4: 

g + gCMB,IR e+ + e- Photons above 1 TeV do not reach us from beyond 10-30 Milion light-years, due to their interaction with IR diffuse background light. Photons above 103 TeV energy do not reach us from the edge of our galaxy because of interaction with the microwave background.

Cosmic Ray spectrum: 

Cosmic Ray spectrum Atmospheric neutrinos Extragalactic flux sets scale for many accelerator models

CR spectrum: 

CR spectrum GZK cutoff or not  ?

Slide7: 

[Slides: Waxman 04]

Slide8: 

Acceleration to 1021eV? ~102 Joules ~ 0.01 MGUT dense regions with exceptional gravitational force creating relativistic flows of charged particles, e.g. coalescing black holes/neutron stars dense cores of exploding stars supermassive black holes

CR acceleration: 

CR acceleration

For a few seconds, a GRB dominates the gamma-ray brightness of the entire Universe: 

For a few seconds, a GRB dominates the gamma-ray brightness of the entire Universe Fig. Credit: Tyce DeYoung

Slide13: 

[Waxman 95] [Frail et al 00]

GZK Sources: 

GZK Sources Sources: GRB √ ; AGN.... #? Rate: RGRB (z=0)~ 0.5 Gpc-3 yr-1 ~ 0.5 10-3 (D/100 Mpc)-3 yr-1 But, arrival time dispersion: tdis~3 107yr (B/10-9 G)2 (B/10 Mpc) (D/100MPC)2 (Ep/1020 eV)-2 NGRB(>Ep, <D) ~ R. tdisp ~104 B-92 B10 D1002 Ep202 GZK event rate: ~ 1 /Km2 /100 yr)  [Waxman 95]

CR data vs. model: 

CR data vs. model [slide: Waxman 05]

Slide19: 

(see talk by S. Coutu)

Slide20: 

Next : how and where are neutrino beams made ?? [ Halzen 02 ] Lab

Slide21: 

radiation enveloping black hole black hole

Sources:: 

Irrespective of the cosmic-ray sources, some fraction will produce pions (and neutrinos) as they escape from the acceleration site through hadronic collisions with gas through photoproduction with ambient photons Cosmic rays interact with interstellar light/matter even if they escape the source Transparent: protons (EeV cosmic-rays) ~ photons (TeV point sources) ~neutrinos Obscured sources Hidden sources Unlike gammas, neutrinos provide unambiguous evidence for cosmic ray acceleration! Sources:

Detection principle: 

Detection principle p  + N +X   Neutrino passes through the earth Interacts near the detector to produce a muon  Cherenkov light is observed

Cherenkov light from muons and cascades: 

Cherenkov light from muons and cascades Maximum likelihood method Use expected time profiles of photon flight times cascade Reconstruction muon

Slide27: 

Nevents ~ Pn -->Area Time  Neutrino flux required to observe N events: 5x10-12 Area (km2) Time (yr) erg cm2/s ntarget  Range ~ 10-4 for 100 TeV neutrinos Detection Probability:   Nevents  =

Slide28: 

Aerial view of South Pole AMANDA Dome Skiway (for planes!) South Pole 1 km IceCube

Building AMANDA: The Optical Module and the String: 

Building AMANDA: The Optical Module and the String

The AMANDA Detector: 

Hot-water-drill 2km-deep holes & insert strings of PMTs in pressure vessels. AMANDA-B10: 302 PMTs, completed in 1997 Old & new A-B10 results presented AMANDA-II: 677 PMTs, completed in 2000 Prelimin. results presented AMANDA challenges: Natural medium Remote location Unfettered bkgd. source Prototype detector AMANDA-II The AMANDA Detector

IceCube: 

IceCube 80 Strings 4800 PMT Instrumented volume: 1 km3 (1 Gton) IceCube is designed to detect neutrinos of all flavors at energies from 107 eV (SN) to 1020 eV

Slide32: 

even neutrons do not escape neutrons escape neutrinos associated with the source of the cosmic rays?

Neutrino ID (solid)Energy and angle (shaded): 

Neutrino ID (solid) Energy and angle (shaded) Neutrino flavor Filled area: particle id, direction, energy Shaded area: energy only

KM3NeT: 

KM3NeT Km3 water Cherenkov detector Deployment approx. 2010 Complement ICECUBE: sc,abs~(100,10) H20, sc,abs~(20,100) Ice Northern site: at lower E , complementary sky coverage EU collaboration Site :Mediterranean Sea based on: NESTOR, NEMO, ANTARES

p + ,p → UHE , : 

p + ,p → UHE , If protons present in (baryonic) jet → p+ Fermi accelerated (as are e-) p, → → ,→e,,e, (-res.: Ep E  0.3 GeV2 in jet frame) → E,br  1014 eV for MeV s (int. shock) → E,br  1018 eV for 100 eV s (ext. rev. sh.)  ICECUBE →0 →2 →  cascade  GLAST, ACTs.. (Waxman-Bahcall 1997;99; Boettcher-Dermer 1998; 00; ) Test hadronic content of jets (are they pure MHD/e , or baryonic …?) Test acceleration physics (injection effic., e, B..) Test scattering length (magnetic inhomog. scale?..or non-Fermi?..) Test shock radius:  cascade cut-off:  ~ GeV (internal shock) ;  ~ TeV (ext shock/IGM) → photon cut-off: diagnostic for int. vs. ext-rev shock

UHE  in GRB 6 possible collapsar GRB -sites: 

UHE  in GRB 6 possible collapsar GRB -sites 1) at collapse, make GW + thermal s (MeV) 2)  If jet outflow is baryonic, have p,n → p,n relative drift, pp/pn collisions → inelastic nuclear collisions → VHE  (GeV) 3 Int. shocks while jet is inside  can accel. protons → p, pp/pn collisions → UHE  (TeV) 4 Int. shocks outside  accel. protons → p collisions → UHE  (100 TeV) 5) ← Ext. rev. shock → EeV  (1018 eV) 6) If supranova shell present outside (SN ocurred >2 days before GRB?) → p, pp of jet protons on shell targets → UHE  (> TeV) [..now constrained] e- capt p,n p, pp p 1,2 3,4 5,6

“Hadronic” GRB Fireballs:Thermal p,n decoupling → VHE ,: 

“Hadronic” GRB Fireballs: Thermal p,n decoupling → VHE , p,n in fireball move together while tpn < texp (rad. acts on p, elastic scatt. couples p,n) p,n decouple when tpn >texp , where pn1, vrel c, pn inelastic; this occurs for  >  ~400 (Derishev etal 99; Bahcall,Meszaros 00; Fuller etal 00) Inelastic pn  ±,e±,e , → 0 → 2  : 5-10 GeV → ICECUBE? det @ z1, R7/yr from all GRB, but only if larger PMT density -rays: 02 , → GLAST,  10 GeV, detect @ z < 0.1 (Bahcall & Meszaros 2000 PRL 85:1362); Lemoine 2002; Beloborodov, 2002 p n p n rdec

While jet is inside progenitor:: 

While jet is inside progenitor:

(2) Jet inside star: GRB , Precursor: 

(2) Jet inside star: GRB , Precursor Jet propagating through progenitor, BEFORE emerging from stellar envelope, can have int. shocks which accel. p+ → p on unobserved X-rays , → ±,  pp, pn on stellar envelope → ±,    few TeV neutrino precursor If progenitor has H-layer R1012 cm (BSG) → Rate(  , TeV ) prec > Rate(  , 100 TeV )int.shock ( easier to detect in ICECUBE ) but, if WR (He core), R1011 cm → Rate(  , TeV) prec < Rate(  , 100 TeV ) int.shock → test progen. size (e.g. @ high z : popIII?) If jet DOES NOT escape ⇒ “choked” jet, s escape, s don’t → “hidden  source” If jet break-out: → photon flashes (3)  Blue - spectrum: ~100 TeV p,→ from shocks outside star Meszaros , Waxman 01 PRL 17 1102 Razzaque, Mészáros, Waxman 03 PRD 68, 3OO1) WR H

When jet is outside progenitor star: GRB internal & external shocks: 

When jet is outside progenitor star: GRB internal & external shocks

s from p in internal & external shocks in GRB: 

s from p in internal & external shocks in GRB Shocks accelerate p+ (as well as the e- which produce MeV ) -res.: E’p E’ 0.3GeV2 in comoving frame, in lab: → Ep ≥ 3x106 Γ22 GeV → E ≥ 1.5x102 Γ22 TeV Internal shock p,MeV → ~100 TeV s External shock p,UV → ~ 0.1-1 EeV s Diffuse flux: detect in km3 Waxman, Bahcall 97 PRL

GRB 030329: precursor(& pre-SN shell?) with ICECUBE : 

GRB 030329: precursor (& pre-SN shell?) with ICECUBE Razzaque, Mészáros, Waxman 03 PRD 69, 23001 Burst of L1051 erg/s, ESN 1052.5 erg, @ z0.17, 68o Prob.of  interaction Flux of 

Core collapse SN : slow jets? : 

Core collapse SN : slow jets? Maybe all core coll. (or Ib/c) SN resemble (watered-down) GRB? Evidence for asymmetric expansion of c.c. (Ib/c) SNR: slow jets Γ~ few ? If so, accel protons while jet inside star, p→πμ→ μ (TeV) Diffuse flux: might be interesting (if 100% SNII make jets), but, more interestingly: individual SN in nearby (2-3 Mpc) gals, e.g. M82, NGC253,  detectable (if have slow jets), at a rate ~ 1 SN/few yr, fluence ~ 100 up-muons/SN, negligible background, in km3 detectors - ICECUBE, KM3NeT Razzaque, Mészáros, Waxman ‘04, PRL 93, 181101; (err: ‘05, PRL 94, 9903) Ando, Beacom (Kaons from pp - astro-ph/0502521)

Diffuse UHE  from pop.III collapse: 

Diffuse UHE  from pop.III collapse At z~5-30(?) pop.III , M ~ 30-300 M , core coll →BH+ accr. Buried jets→p→  , → -bursts (but: dep. on stellar rot.rate) Eiso~1054-1056 (?) erg (dep. on BH mass, dM/dt) Detect high z star formation, primordial IMF Recent (8/04) : can constrain w. AMANDA latest results: → Eiso~1056 erg only for ≤1%, → Eiso≥1054 erg for ≤ 50% ! Schneider, Guetta, Ferrara aph/0201342 Recent AMANDA u.l.

AGN as UHE  sources: 

AGN as UHE  sources Big brother of GRB: massive BH (107-108 Msun ) fed by an accretion disk → jet – But, jet jet,agn ~10-20 (while grb ~ 102-103 ) UV photons from disk; in addition, line clouds provide extra photons (+back-scatter) Typical (“leptonic”) model: SSC (sync-self-compton); SEC(sync-exter.compton)

Radio-loud blazars (jet nearly head-on):Mrk 501: 

Radio-loud blazars (jet nearly head-on): Mrk 501 1997 flare: TeV; (GeV: upper lim only w EGRET) GeV detected sometime @ quiescence ←Typical “astrophysical” SSC or ESC “leptonic” jet  model fit But: competing “hadronic” jet  model fits 

Radio-loud hadronic Blazar models (PSB-proton synchrotron blazar - -ray spectrum from cascades): 

Radio-loud hadronic Blazar models (PSB-proton synchrotron blazar - -ray spectrum from cascades) Full : synchrotron  SED (target photons) Dash: p-sync. casc.; Dash-3 dot: ±-sync. casc; Dots: 0 casc; Dash-dot: ± casc (Muecke, et al, Apph, astro-ph/0206164 )

Mrk 501 : protypical HBL : 

Mrk 501 : protypical HBL a)  PSB: Quiet state  b)  PSB: Flare state   e-sync  targets + p-sync  + p, casacdes,  casacdes & sync (Muecke et al, a-ph/0206164) c) → LEP:Flare state  → e-sync  + e-Inv. Compton scatt (Ghisellini et al, e.g. A&A 386, 833 (2002) etc – “standard” astrophysical. picture

Radio-quiet (core) AGN s: 

Radio-quiet (core) AGN s AGN are powered by accretion on massive (106-108 M ) BHs 90% of AGNs are radio-quiet (no jets), core X-ray Core emission model: aborted jet → cloud collisions →shocks → p accel → p →   Diffuse flux: already constrained by latest AMANDA results Alvarez-Muñiz & Mészáros, 2004, PRD 70, 123001

Other Implications of GRB UHE : 

Other Implications of GRB UHE  Special relativity: simultaneity of arrival of , tested to t < 1 s (10-3 s in short bursts) Time delay due to i mass: t (i ) ~ 10-12 (D/100Mpc)(Ei/100TeV)-2(mi /eV)2 s (whereas for SN 1987a t (i )~ 10-8 s ) Vacuum oscillations: at source exp. N ~ 2Ne at observer exp.  ratios ,  upgoing  appearance → sensitive to m2 <10-16 (E/100TeV)(100Mpc/D) eV2 (for m 0.1 eV due to finite pion life mixing is caused by decoherence rather than oscillation)

Slide53: 

GZK neutrinos: GZKbrick wall n GZK neutrinos  probably 2nd most likely source of UHE neutrinos! they are a guaranteed source!! Two predictions 1. There is a brick wall for the highest energy cosmic rays. We should observe energies below about 1020 eV. 2. Reactions that limit the cosmic ray energies produce neutrinos as a by-product Ultra-high energy cosmic rays: Origin; unknown, but Standard Model: Ordinary charged particles accelerated by distant sources: AGN, GRBs… If so: GZK neutrinos are the signature Probably necessary and sufficient to confirm standard GZK model (Courtesy A. Silvestri via D. Saltzberg)

What do we need for a GZK n detector?: 

What do we need for a GZK n detector? Askaryan process: coherent radio Cherenkov emission: EM cascades produce a charge asymmetry  radio pulse Process is coherent  Quadratic rise of power with cascade energy Neutrinos can shower in radio-transparent media: air, ice, rock salt, etc. RF economy of scale very competitive for giant detectors How can we get the ~100-1000 km3 sr yr exposures needed to detect GZK neutrinos at an acceptable rate? ANSWER Standard model GZK: Fn: <1 per km2 per day Only 1 in 500 interact in ice QUESTION: Both AMANDA-II or IceCube may expect to see 1 event every 2 years in its fiducial volume— requires astronomical level of patience!

from GZK CRs to GZK s : 

from GZK CRs to GZK s Seckel & Stanev astroph/050244 2 ≠CR models  same GZK fit  ≠ GZK  flux

Slide56: 

ANtarctic Impulsive Transient Antenna NASA funding started 2003 for full launch in 2006 ANITA-lite succesfully launched & tested Dec 2003 600 km radius, 1.1 million km2

Slide57: 

n ANITA concept Antarctic Ice at f<1GHz, T<-30C Nearly Lossless RF transmission Negligible scattering largest homogenous, RF-transmissive solid mass in the world!

EUSO: 

EUSO ISS project ESA/NASA/RSA/JSA; precursor for OWL (free-flyer) 5.1019 – 1021 eV EECRs, EENUs Monocular 2.5m Fresnel lens, measure EAS through atm. fluor Thresh: 3.1019 eV; Effic. @ 1020 eV : 300-1000 event/yr Launch: 2010-12, but: shuttle ? Possibly: JSA unmanned shuttle

CR &  bounds: 

CR &  bounds

Summary: UHE : 

Summary: UHE  GRBs  1020 eV protons - Predictions 103 km2 area detectors - Experiments: HiRes, Auger, EUSO/OWL GRBs  10 GeV, 1 Tev, 100 TeV, 1018 eV ’s - Flux  1 Gton detectors - Exp’s: AMANDA, IceCube, Antares, Nestor, Nemo  detection - GRBs: CR puzzle, GRB progenitors & physics -  physics:     appearance - Cross sections at E > LHC