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TAIGA experimnent: Ultra-high energy gamma-ray astronomy at Tunka Valley

TAIGA experimnent: Ultra-high energy gamma-ray astronomy at Tunka Valley. Leonid Kuzmichev Skobeltsyn Institute of Nuclear Physics MSU July 201 5 , B . Koty. Gamma-radiation > 0.1 MeV High energy gamma-ray astronomy > 1 GeV

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TAIGA experimnent: Ultra-high energy gamma-ray astronomy at Tunka Valley

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  1. TAIGA experimnent: Ultra-high energy gamma-ray astronomyat Tunka Valley Leonid Kuzmichev Skobeltsyn Institute of Nuclear Physics MSU July 2015, B.Koty Gamma-radiation > 0.1 MeV High energy gamma-ray astronomy> 1 GeV Very high energy gamma-ray astronomy(VHE)> 100 GeV Ultra high ebergy gamma-ray astronomy(UHE) > 10 ТeV

  2. 2lectures • High-energy gamma-ray astronomy - introduction • 2. Gamma-ray observatory TAIGA • «Physics» of gammaа-ray astronomy • Origin of Cosmic rays • 2. Intrinsic structure of astrophysical objects • 2 Gamma-astronomy andcosmology • 3. Dark matter

  3. Lecture 1:High energy gamma-ray astronomy

  4. Plan 1.Introductions – aims of gamma-rays astronomy 2. Gamma-rays generation Synchrotron radiation Inverse Compton scattering Pi-0 decay 3. Absorption of gamma-rays 4. Gamma-rays and Super Nova Remnants (SNR) 5. Status of arrays and projects

  5. How we study Universe OK 1 TeV = 1012 eV 1 PeV = 1015 eV 1 EeV = 1018 eV 1 ZeV = 1021 eV

  6. Gamma-ray astronomy very successful part of astrophysics • 150 TeVsources • 2000 GeV sources • Instruments operation: • VERITAS • HESS • MAGIC • Fermi-Lat • Argo-YbJ • Tibet-III • 7. HAWC • Projects: • CTA - 2017-18 ? • LHAASO - 2017-2018( ?) • 3. TAIGA

  7. TeV hi PeV

  8. More than 100 local sources of gamma rays with energy >100 GeV

  9. For this energy range (> 30 TeV) (1-10 )km2 area arrays are needed

  10. Galactic gamma-rays sources

  11. 2. Gamma-rays generation • 1.Synchrotron radiation • 2. Inverse Compton scattering • 3. Pi-0 decay

  12. Synchrotron radiation νLarmor = 2.8 106 ( B/ 1 G) Гц νсинх.=3/2 νLarmor ( E / m c2 )3- X-rays Intensity ~ B2 flux of electrons

  13. Inverse Compton –effect on relict photons Cross sections: σ = σthom. ( S < (mc2)2 ) σ = σthom(mc2)2 / s ( S > (mc2)2 σтhom= 0.66 10-25 cm2 εγ ~10-4 -10-3 эВ Eγ ~εγ х ( E / mc2 )2 100 ТeV electron20 ТeVphotons

  14. Relation between fluxes and energy Photon from Inverse Compton -effect E γ = 2 ( εx / 0.1 keV) (B/10 µG)-1 Synchrotron radiation 1 erg ~ 1 TeV Flux of energy: f (E) = E2 dN/dE ( erg/cm2 sec) (f (Eγ)) inverseCompton f(εx) synchrotron = 0.1 (B/10 µG)-2

  15. Gamma-rays frompi-decay P + P π0 + All E γ~ 0.1 E p 2 γ Energy spectrum Energy spectrum of protons :A E –γ Energy spectrum of gamma-rays : В E-γ why?

  16. Synchroyton radiation Inverse Compton dNe / dE ~ E –αdNγ/dE ~ E-(α+1)/2

  17. Gamma ray fromprotons dN/dE ~ E-p - pi-0 - gamma-raysdN/dE ~ E -p From electrons dN/dE ~ E-p - gamma-raysdN/dE ~ E –(p+1).2

  18. Absorption ofgamma-rays γ + photon → e+ + e- University λ (Mpc) Galaxy Exp ( - l / λ) proton 1 Пэв 8 kpc 40 kpc

  19. Distance from the nearest galaxies LMC - 160 kpc SNR 1987 Androdema 2 Mpc M82 ( startburst galaxy) 4 Mpc Markarian 421 120 Mpc ( the nearest Blazar)

  20. 3.Gamma-astronomy and supernova remnants

  21. SNRs – main sourses of Galactic Cosmic rays 1.1933 – Baade and Zwicky– Explosion of SN – source of CR 2. 1949 – Fermi – first theory ofcosmic rays acceleration 3. 1963 – Ginzburg, Sirovatsky – transfer of 10% of kineticenergyof SNR КЛ is enough to explain intensity of CR 4. 1977 – 1978 -Krymsky, Bellat all – theory of acceleration on shock waves 5. 1993-1996 – Berezhkoet al. – nonlineralytheory acceleration. 6. 2003-2005 – Bell, Berezhko &Volk, Ptuskin &Zirakoshvily – amplification of magnetic field on the front ofshockwaves – Emax ~ Z · 1015 eV Cas A radio polarization in red (VLA), X-rays in green (CHANDRA), optical in blue (HST) SN explosion – 1053 erg Kinetic energy of - 1051 erg Rate- 1 per 30 y

  22. Observations nonthermal X-rays εkeV = 1 BμG(Ee/120 TeV)2 εmax ~ 100 TeV radio emission νMHz = 4.6 BμG (Ee,GeV )2 E = 50 MeV – 30 GeV (100 GeV for IR) γ = 1.9 – 2.5 We = 1048 – 1049 erg Ginzburg & Syrovatskii 1964 Shklovsky 1976 synchrotron γ e SNR inverse Compton εγ = ε0(Ee/mec2)2 γ p TeV γ– rays electrons/protons εmax ~ 100 TeV e π0 γ

  23. Cosmic Ray diffusive acceleration in Supernova Remnants ESN ~ 1051erg Krymsky 1977 Bell 1978 =hock compression ratio - is the shock compression ratio for strong shock Frequent scatterings of CRs on magnetic field irregularities in collisionless medium provide efficient acceleration of CRs at strong shock

  24. Maximum Energy Emax = U x R x B - Hillas rule Amplification of magnetic field By stream instability

  25. Magnetic field amplification –steaming instability ρISM Beff Results of modeling (Lucek & Bell, 2000) & theoretical considerations (Bell 2004; Pelletier et al. 2006; Zirakashvili et al. 2007)+ Spectral properties of SNR synchrotron emission + Fine structure of nonthermal X-ray emission VS BISM CR flux SNR magnetic field is considerably amplified L shock B ~ 200 n1/2 ( U/ 104 km/s ) 3/2 µG Beff2/8π ≈ 10-2ρISMVS2 Beff >> BISM Emax ~ 200 n1/2 ( U/ 104 km/s ) 2 Rpc TeV

  26. Chandra SN 1006 Chandra Cassiopeia A • Filamentary structure of X-ray emission • of young SNRs • consequence of strongly amplified magnetic field, • leading to strong synchrotron losses

  27. Dependence of Emax from time Emax ~ U R B ~ U2 R R = R0 ( t/t0) 2/5 Sedov solution t0 – beginning of Sedov phase t0 ~ n -1/3 M5/6 E1/2 - 100 -1000 years Emax ~ t -4/5

  28. «Cooling» ofelectrons T 1/2 - time of transfer energy from eletrons to gamma-rays dE/dt = b E2 W CMB = 0.25 eV/cm3 b = 4/3 (σT c )/ (mc2 )2 ( W CMB + B2/ 8π ) σT h = 6.6 10-25 cm2 Sychrotron radiation Hight-energy gamma-rays E (t ) = Eo/ ( 1 +bt E0) T1/2 = 1/bE0 For E0 =20 TeV T = 5 10 4 y W CMB / W B = 0.1 (B/10 µG )-2

  29. «Cooling» ofprotons : Ƭpp = 1/ ( ngas · c · k · σpp) = 6 x 107 ( n gas / 1 cm-3)-1год, k - inelasticity

  30. Relation between fluxesof-gamma-rays Fγ ( IC) = W e / T F γ (π ) = Wp /Ƭ Fγ ( IC) / F γ (π ) = 10 3 ( We/ Wp) (n/1 cm-3)-1 forEγ =1 TeV

  31. Derivation of DAV formula : Ƭpp = 1/ ( ngas · c · k · σpp) = 6 x 107 ( n gas / 1 cm-3) год, k - inelasticity P (E) = I x E-2 total energy = integralfrom 1 GeVto 1 PeV I x ln( 10^6) = 1050 erg I = 7x 10^48 erg c (P -γ) = 0.1 - part of proton energy transfered to gamma-ray Fγ (Eγ) E ^2 = I x c (P -γ) / (Ƭpp ( 4π d^2) ) = 10^(-11)  ( Wcr/ 10^50 erg) ( n gas / 1 cm-3 )  ( d/ 1kpc)^-2 erg / cm2 s

  32. SNR detected at TeV energies

  33. Molecular clouds : possibility tocatchPeVatrons If clouds is in a distance of 100 pc from SNR flux will be in 10 times smaller, but duration Will be in 20 time longer – near to 10000 y/ Mass (104 - 105 ) M of Sun Density ~100 g/ cm2 1% of Galaxy volume mostly from H2

  34. Gamma-ray astronomyand Crab nebular • Explosion in 1054 , distance 2 kp, • In the centre of nebular –pulsar with 33 ms period • First reliably registrated gamma-ray source • T.Weeks 1989. 9 RMS • Steady gamma-radiation – standard candle • “pulsate”gamma-radiation • Gamma «bursts»

  35. Inserlude: extensiveatmospheric showers (EAS)

  36. P, A Xmax – maximum of EASdevelompent increasingnumber of particles Хmax = A + B Ln( E/A) 20-30 km Atomic number Nmax  E number of particles Energy 3-5 км Half of the particle in the circle of 80 m Shower particles Shower core Energy of particles: Electrons: 30-100 MeV Muons 0.5 Gev Particle detectors

  37. P, A Registration of Chrenkov light forEe >25 MeV Ve > C/ n – speed of lightin the air Cos (tet) =1/n tet ~0.5 deg 20-30 km Cherenkov light Q tot  E Phononsdetectors

  38. Energy threshold of Cherenkov array Cherenkov pulse Sd– areaof PMT QE- quantum efficiency Pph– поток черенковских фотонов T - duration of pulse ( 10 - 20 ns)  - FOV Iф – night light background  3.1012 T Ss•Pph•QE signal 5 = noise  •Sд •Iф•T Pph ~ E - energy of EAS  Iф•  •T Ethersh~ m-2 sec1  Sд •  Sd~ 0.1 m2 иQE0.1 : Eth 200 ТeV

  39. How select gamma-shower from proton shower F( gamma) x S x T Signal (background )1.2 = = 5 (F(CR)x S x Ω x T)1/2 Background selection : Imagingatmospheric Cherenkov telescope poor muons showers ( in gamma showers in 30 times smaller muons than inproton shower) K signal, Kback –rejection factors. Q = K signal / sqrt ( K background)

  40. Short history of gamma-rays astronomy

  41. Cherenkov Technique used for Gamma Ray Astronomy Crimea Experiment 1959-1965, Chudakov, et al., (SNR, radio galaxies)

  42. First Gamma-ray Experiment at Whipple Observatory, 1967-68 The pioneer, the #1 in gamma astronomy

  43. The Pioneer Trevor Weekes and his 10m Ø Whipple telescope gave birth to g-ray astrophysics: 9s from Crab Nebula in 1988 ! „If a telescope can within a few s evaporate a solid piece of steel, it can also measure gamma rays“ ;-)

  44. Instruments

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