TAIGA experimnent Ultrahigh energy gammaray astronomy at Tunka

TAIGA experimnent: Ultra-high energy gamma-ray astronomy at Tunka Valley Leonid Kuzmichev Skobeltsyn Institute of Nuclear Physics MSU July 2015, B. Koty Gamma-radiation > 0. 1 Me. V High energy gamma-ray astronomy > 1 Ge. V Very high energy gamma-ray astronomy (VHE)> 100 Ge. V Ultra high ebergy gamma-ray astronomy (UHE) > 10 Тe. V

2 lectures 1. High-energy gamma-ray astronomy - introduction 2. Gamma-ray observatory TAIGA «Physics» of gammaа-ray astronomy 1. Origin of Cosmic rays 2. Intrinsic structure of astrophysical objects 2 Gamma-astronomy and cosmology 3. Dark matter

Lecture 1: High energy gamma-ray astronomy

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

How we study Universe OK 1 Te. V = 1012 e. V 1 Pe. V = 1015 e. V 1 Ee. V = 1018 e. V 1 Ze. V = 1021 e. V

Gamma-ray astronomy very successful part of astrophysics 150 Te. V sources 2000 Ge. V sources Instruments operation: 1. VERITAS 2. HESS 3. MAGIC 4. Fermi-Lat 5. Argo-Yb. J 6. Tibet-III 7. HAWC Projects: 1. CTA - 2017 -18 ? 2. LHAASO - 2017 -2018( ? ) 3. TAIGA

Te. V Pe. V hi

More than 100 local sources of gamma rays with energy >100 Ge. V

For this energy range (> 30 Te. V) (1 -10 )km 2 area arrays are needed

Galactic gamma-rays sources

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

Synchrotron radiation ν Larmor = 2. 8 106 ( B/ 1 G) Гц νсинх. = 3/2 ν Larmor ( E / m c 2 )3 - X-rays Intensity ~ B 2 flux of electrons

Inverse Compton –effect on relict photons Cross sections: σ = σthom. σтhom= 0. 66 10 -25 cm 2 ( S < (mc 2)2 ) σ = σthom (mc 2)2 / s ( S > (mc 2)2 εγ ~10 -4 -10 -3 э. В Eγ ~ εγ х ( E / mc 2 )2 100 Тe. V electron 20 Тe. V photons

Relation between fluxes and energy Photon from Inverse Compton -effect E γ = 2 ( εx / 0. 1 ke. V) (B/10 µG)-1 Synchrotron radiation 1 erg ~ 1 Te. V Flux of energy: f (E) = E 2 d. N/d. E ( erg/cm 2 sec) (f (Eγ)) inverse Compton f(εx) synchrotron = 0. 1 (B/10 µG)-2

Gamma-rays from pi- decay P + P π0 + All 2 γ E γ ~ 0. 1 E p Energy spectrum of protons : A E –γ Energy spectrum of gamma-rays : В E- γ why?

Synchroyton radiation Inverse Compton d. Ne / d. E ~ E –α d. Nγ /d. E ~ E-(α+1)/2

Gamma ray from protons d. N/d. E ~ E-p - pi-0 - gamma-rays d. N/d. E ~ E -p From electrons d. N/d. E ~ E-p - gamma-rays d. N/d. E ~ E –(p+1). 2

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

Distance from the nearest galaxies LMC - 160 kpc SNR 1987 Androdema 2 Mpc M 82 ( startburst galaxy) 4 Mpc Markarian 421 120 Mpc ( the nearest Blazar)

3. Gamma-astronomy and supernova remnants

SNRs – main sourses of Galactic Cosmic rays 1. 1933 – Baade and Zwicky– Explosion of SN – source of CR 2. 1949 – Fermi – first theory of cosmic rays acceleration Cas A 3. 1963 – Ginzburg, Sirovatsky – transfer of 10% of kinetic energy of SNR КЛ is enough to explain intensity of CR 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 4. 1977 – 1978 -Krymsky, Bell at all – theory of acceleration on shock waves 5. 1993 -1996 – Berezhko et al. – nonlineraly theory acceleration. 6. 2003 -2005 – Bell, Berezhko & Volk, Ptuskin & Zirakoshvily – amplification of magnetic field on the front of shock waves – Emax ~ Z · 1015 e. V

Observations nonthermal X-rays radio emission νMHz = 4. 6 BμG (Ee, Ge. V )2 E = 50 Me. V – 30 Ge. V (100 Ge. V for IR) γ = 1. 9 – 2. 5 We = 1048 – 1049 erg Ginzburg & Syrovatskii 1964 Shklovsky 1976 synchrotron γ e SNR π0 γ εke. V = 1 BμG(Ee/120 Te. V)2 εmax ~ 100 Te. V p e Compton γ inverse ε = ε (E /m c 2)2 γ 0 e e Te. V γ – rays electrons/protons εmax ~ 100 Te. V

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

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

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)+ VS BISM Spectral properties of SNR synchrotron emission + Fine structure of nonthermal X-ray emission CR flux shock SNR magnetic field is considerably amplified L 2 -2 Beff /8π ≈ 10 ρISMVS Beff >> BISM 2 B ~ 200 n 1/2 ( U/ 104 km/s ) 3/2 µG Emax ~ 200 n 1/2 ( U/ 104 km/s ) 2 Rpc Te. V

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

Dependence of Emax from time Emax ~ U R B ~ U 2 R R = R 0 ( t/t 0) 2/5 Sedov solution t 0 – beginning of Sedov phase t 0 ~ n -1/3 M 5/6 E 1/2 - 100 -1000 years Emax ~ t -4/5

«Cooling» of electrons T 1/2 - time of transfer energy from eletrons to gamma-rays d. E/dt = b E 2 W CMB = 0. 25 e. V/cm 3 b = 4/3 (σT c )/ (mc 2 )2 ( W CMB + B 2/ 8π ) E (t ) = Eo/ ( 1 +bt E 0) Hight-energy gamma-rays σT h = 6. 6 10 -25 cm 2 Sychrotron radiation T 1/2 = 1/b. E 0 For E 0 =20 Te. V T = 5 10 4 y W CMB / W B = 0. 1 (B/10 µG )-2

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

Relation between fluxes of-gammarays Fγ ( IC) = W e / T F γ (π ) = Wp /Ƭ Fγ ( IC) / F γ (π ) = 10 3 ( We/ Wp) (n/1 cm-3)-1 for Eγ =1 Te. V

Derivation of DAV formula : Ƭpp = 1/ ( ngas · c · k - inelasticity σpp) = 6 x 107 ( n gas / 1 cm-3) год, P (E) = I x E-2 total energy = integral from 1 Ge. V to 1 Pe. V I x ln( 10^6) = 1050 erg I = 7 x 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) ( ( d/ 1 kpc)^-2 erg / cm 2 s n gas / 1 cm-3 )

SNR detected at Te. V energies Name RX J 1713. 7 Dist(kpc) 1 Size, pc Age, y 17. 4 L/ 1033 erg/c 1600 8 Ѓ (slope) 2. 0 RX J 0852 Vela Junor 0. 2 (1) 6. 8(34) 400(5000) 0. 26(6. 40) 2. 2 RCW 86 1(2. 50) 11(28) 1600(10000) 1(6) 2. 5 SN 1006 2. 2 18. 3 1000 1. 24 2. 3 Cas A 3. 4 2. 5 350 7 2. 4 Tycho 1572 3. 5 6 443 0. 1 1. 95 SNR G 353. 6 -07 3. 2 27 2500 (14000) 10 2. 3

Molecular clouds : possibility to catch Pe. Vatrons 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/ cm 2 1% of Galaxy volume mostly from H 2

Gamma-ray astronomy and 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 1. Steady gamma-radiation – standard candle 2. “pulsate” gamma-radiation 3. Gamma «bursts»

Inserlude: extensive atmospheric showers (EAS)

P, A increasing number of particles 20 -30 km Хmax = A + B Ln( E/A) Nmax 3 -5 км Shower particles Xmax – maximum of EAS develompent number of particles E Atomic number Energy Half of the particle in the circle of 80 m Shower core Energy of particles: Electrons: 30 -100 Me. V Particle detectors Muons 0. 5 Gev

P, A Registration of Chrenkov light for Ee >25 Me. V Ve > C/ n the air – speed of light in Cos (tet) =1/n 20 -30 km tet ~0. 5 deg Cherenkov light Q tot E Phonons detectors

Energy threshold of Cherenkov array Cherenkov pulse T signal noise = Sd– area of PMT QE- quantum efficiency Ss • Pph • QE • Sд • Iф • T Pph ~ E - energy of EAS Ethersh ~ Iф • • T 5 Pph– поток черенковских фотонов T - duration of pulse ( 10 - 20 ns) - FOV Iф – night light background Sд • Sd ~ 0. 1 m 2 и QE 0. 1 : Eth 200 Тe. V 3. 1012 m-2 sec 1

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

Short history of gamma-rays astronomy

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

First Gamma-ray Experiment at Whipple Observatory, 1967 -68 The pioneer, the #1 in gamma astronomy 56

The Pioneer Trevor Weekes and his 10 m Ø Whipple telescope gave birth to g-ray astrophysics: 9 s 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“ ; -) 57

Instruments

Number of muons in proton EAS in 30 times more than in gamma-EAS

HAWC (High Altitude Water Cherenkov) Te. V Gamma-Ray Observatory

HAWC Site Location in Mexico • 4100 m (13, 500’) above sea level • Latitude of 19 deg N • Temperate Climate • Existing Infrastructure HAWC Large Millimeter Telescope (50 m dia. dish) Pico de Orizaba 5600 m (18, 500’)

NEW projects

An observatory for ground based gamma-ray astronomy: CTA 3 types of mirrors: 24 m diameter , FOV 4 -5 deg 10 m , FOV 6 -8 deg ( 100 Ge. V – 10 Te. V) 4 -6 m FOV 10 deg > 10 Тe. V

LHAASO——The third generation of survey facility for VHE γray sources 24 Wide FOV air Cherenkov image Telescopes. For Air Fluorescence measurement above 0. 1 Ee. V also 400 burst detectors For high energy Secondary particles Near the core of air showers 6100 scintillator detectors and 1200 μdetector s form an array covering 1 km 2 LHAASO Layout in 1 km 2 at 4300 m a. s. l. 90 k square meter Water Cherenkov Deter array. Each one has a size of HAWC

Survey for γ-sources very detailed spectroscopy investigation ( a few hundred extra-galactic sources are expected ) 1000 10 Ev/yr

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