Doctoral Thesis. Detection of TeV gamma-rays from the Supernova Remnant RX J

Transcription

1 Doctoral Thesis Detection of TeV gamma-rays from the Supernova Remnant RX J Hideaki Katagiri Department of Physics, Graduate School of Science The University of Tokyo Hongo, Bunkyoku, Tokyo , Japan December 19, 2003

2 Abstract Sub-TeV gamma-rays emitted from the northwest rim of the supernova remnant RX J , where maximum non-thermal X-rays were detected by ASCA, were observed by the CANGAROO-II 10-m imaging air Cherenkov telescope (IACT) at South Australia in 2002 and Data obtained in 187 hours of observation time gave the 6.4 σ statistical significance using the image analysis with the likelihood method. The flux of gamma-rays was 0.25 ± 0.06 times that of the Crab nebula at 500GeV with the spectral index of 4.5 ± 0.7 above 300GeV. The α (image orientation angle) distribution indicated a marginally extended emission, but was still consistent with a point source within statistical errors. The center of the obtained morphology coincided with the X- ray maximum point. The gamma-ray spectra were estimated under the assumptions of the synchrotron/inverse Compton model and decay of π 0 s produced by proton-nucleon collisions. Our data strongly favored TeV gamma-ray emission from π 0 decay. A total cosmic-ray energy of to erg is required, when the molecular cloud density is 5000 to 50 protons cm 3, assuming the distance was 0.5 kpc. The two-zone model of synchrotron/inverse Compton model with fine structures of X-ray emissions, however, can also explain the broadband spectrum. Further observations and analyses such as XMM- Newton and Chandra X-ray satellites, NANTEN radio telescope, and CANGAROO-III stereoscopic system (IACTs) are strongly awaited to confirm the emission mechanism.

3 Contents 1 Introduction 8 2 Review Origin of Cosmic Rays Fermi Acceleration Diffusive shock acceleration (DSA) Observations of TeV gamma-rays from Supernova Remnants SNR RX J (G ) Processes of Non-thermal Emissions π 0 decay Synchrotron Radiation Inverse Compton Scattering Bremsstrahlung Imaging Air Cherenkov Technique Overview Extensive Air Showers Electromagnetic Showers Hadronic Showers Cherenkov Radiation Imaging Air Cherenkov Technique The CANGAROO-II 10-m Telescope Reflector Imaging Camera Electronics and Data Acquisition System Observations and Calibrations Observations Calibrations Terminologies Field Flattening Time-walk Corrections Rejection of Bad Channels DST ADC Conversion Factor

4 6 Analysis Reduction of the Night Sky Background (NSB) ADC Distributions Clustering TDC Cut Cloud Cut and Elevation Cut Selection of Bad Pixels due to Starlights and Electrical Noises Selection of Bad Pixels using ADC Distributions Image Analysis Monte-Carlo Simulations Image Analysis using Likelihood Method α Distributions Differential Fluxes Results α Distributions Effective Area and Energy Threshold Differential Fluxes Morphology Various checks Conventional Cut Effects of the Bad Pixel Cut Hillas Parameter Distributions of Excess Events Crab Analysis Signal Rate Systematics Discussion Broadband Spectrum Synchrotron/inverse Compton Model π 0 Decay produced by Proton-nucleon Collisions Summary of Discussions Conclusion 123 A Definitions of the Image Parameters 129 B Differential Cross section 131 2

5 List of Figures 2.1 Balloon flights of Hess Differential energy spectrum of cosmic rays Cosmic ray elemental abundances measured at Earth compared to the solar system abundances, all relative to silicon Spectrum of the ratio between the number of Boron and that of Carbon in cosmic rays as a function of kinetic energy Schematic view of the diffusive shock acceleration around the shock front in the laboratory frame Number of the sources detected with various detectors versus year X-ray of SN1006 obtained by ASCA satellite Contours of statistical significance map of gamma-rays from SN1006 northeast rim obtained by CANGAROO 3.8m telescope (CANGAROO-I) Schematic view of the Spectral Energy Distributions for synchrotron/inverse Compton model and π 0 decay model Spectral Energy Distribution observed from the NE rim of SN Contours of statistical significance map of RX J northeast rim obtained by CANGAROO 10m telescope Spectral Energy Distribution from RX J , and emission models Soft X-ray images of Vela SNR observed by ROSAT Hard X-ray image of RX J observed by ASCA GIS Radio images at 4.85GHz observed by the Parkes radio-telescope, centered on RX J Integrated intensity map of CO obtained with NANTEN 4 m milli-metre radio telescope and a soft X-ray image by ROSAT Spectral distribution of the power of the total (over the directions) radiation from synchrotron radiation Spectrum of bremsstrahlung Spectrum of pair creations Definitions of the variables representing the height and depth of the atmosphere Schematic view of electromagnetic showers of 1TeV gamma-rays in the air Schematic interaction processes of hadronic showers Showers of gamma-rays and protons in 1TeV simulated by the Monte Carlo methods Schematic view of Cherenkov radiation Images and the lateral distributions of photons produced from gamma-ray showers in 1TeV

6 3.9 Images and the lateral distributions of photons produced from proton showers in 3TeV Transmission of the Cherenkov photons Examples of the distributions of photons on the camera plane of IACTs Hillas parameters and α Schematic view of images generated from gamma-rays and protons Distributions of α about Markarian 421 obtained by the Whipple Observatory CANGAROO-II 10-m telescope Schematic illustration of the cross section of a segmented mirror Imaging camera of CANGAROO-II 10-m telescope PMT (Hamamatsu R4124UV) used for the CANGAROO-II 10-m telescope Spectral response of photocathod Light guides of CANGAROO-II 10-m telescope Block diagram of DAQ for the CANGAROO-II 10-m telescope Schematic diagram of the TKO front-end module and the discriminator and summing module Updated discriminator Schematic diagram of the event trigger logic Arrival time (TDC) distributions ADC distributions for pixels, after pedestal subtraction TDC distributions for various cluster sizes Distributions of TDC for a typical run after T5a-clustering and adjusting the mean TDC of each event to Change of event rate due to clouds and the change of elevation Shower rates versus cosine of the zenith angle in 2002 and Distributions of scaler counts for pixels Optical image around the NW rim of RX J taken by Digital Sky Survey Scalar maps with the correction of the rotation of the field of view using the data in Tracks of the stars on the focal plane during the observation around the NW rim of RX J Integral observation time distribution on the focal plane for the bright stars Distribution of χ 2 /DOF of entries and ADC distributions Examples of good pixels and bad pixels in Examples of good pixels and bad pixels in Zenith angle distributions in 2002 and Distributions of Width, Length, Distance, and Asymmetry E ratio distributions Correlations between Hillas parameters (Width and Length) and the logarithm of ADC sum Image parameters of the OFF-source data and the Monte-Carlo simulations of protons Probability Density Functions normalized to unity

7 6.19 Distributions of Likelihood-ratio Figure of merit versus Likelihood-ratio cut and acceptance versus Likelihoodratio cut Distributions of α for OFF-source data in Distributions of α for the OFF-source data and the Monte-Carlo simulations of protons after the likelihood cut Acceptances and the acceptances/ α versus α-cut values obtained by the Monte-Carlo simulations of the gamma-rays with a point-source assumption Distributions of α: 2002, 2003, and the combined Spectra of the Monte-Carlo simulation of gamma-rays under the assumption of the spectrum with the spectral index of 2.5 as a function of energy and ADC sum which is proportional to the energy Effective areas of the gamma-rays under the assumption of the spectrum with the spectral index of 2.5 as a function of energy α distributions for each ADC sum Correlation between the energy and the ADC sum of the gamma-ray events from the Monte-Carlo simulation Distributions of the energies of the gamma-rays in each region of ADC sum Differential fluxes with the statistical errors Assumed indices of the gamma-ray Monte-Carlo simulation and the ratio of the indices between the assumed spectrum and the obtained spectrum Distributions of α after the iteration Distributions of α for OFF-source data for various χ 2 /DOF cuts Distributions of α for the Monte-Carlo simulations assuming the emission is extended and the ratio N(α < 15 )/N(α < 30 ) Spectrum of gamma-rays by the Monte-Carlo simulation assuming a spectral index of 4.5 as a function of energy Effective areas of the gamma-rays under the assumption of the spectrum with the spectral index of 4.5 as a function of energy Distributions of α after the iteration for each ADC sum Differential fluxes with the statistical errors after the iteration Differential fluxes of different binnings Significance maps of gamma-ray signal Acceptance versus offset angle of the gamma-ray source position from the center of the field of view α distributions after the conventional cut α distributions after various image analyses Bad pixels selected using the ADC distributions Changes of the distributions for α, Width, and Length obtained by analyzing the data of the Monte-Carlo simulations of gamma-rays Distributions of the Hillas parameters for the excess events and the gammarays generated by the Monte-Carlo simulations α distribution of Crab nebula Differential flux of Crab nebula Significance map of Crab nebula

8 8.9 Excess events of the gamma-ray signals as a function of the observation time for each run Differential fluxes obtained by the various trigger conditions Distribution of χ 2 /DOF of the entries and the normalized ADC distributions Differential fluxes by the various χ 2 /DOF cuts of the entries and the normalized ADC distributions Variation of differential fluxes obtained by the Monte-Carlo simulations by changing the spectral indices and the extents of emission Differential fluxes obtained by the various L ratio cuts α distributions and the differential fluxes obtained by the various cut values of L ratio Distributions of the differential fluxes obtained from the various conditions of the analysis in each region of ADC sum Distributions of the energies obtained from the various conditions of the analysis in each region of ADC sum Differential fluxes with all errors Distributions of the excess events obtained by the analyses with various assumptions and methods EGRET gamma-ray intensity map near RX J based on archival mapped data Gamma-ray emissivities for π 0 decay, bremsstrahlung, and inverse Compton scattering for various radiation fields One-zone synchrotron/inverse Compton models Ratio between the Klein-Nishina cross section and the cross section of Thomson scattering as a function of the logarithm of the emitted photon energy Synchrotron/inverse Compton models: cross section of the Thomson scattering and Klein-Nishina cross section Synchrotron/inverse Compton models with the different sizes of the emission regions in X-rays and TeV gamma-rays (two-zone model) X-ray images of ASCA GIS based on the archival data Synchrotron/inverse Compton models with the different spectral indexes α and magnetic fields B, and maximum accelerated energy of the electrons E max between X-rays and TeV gamma-rays Allowed region determined from the χ 2 /DOF values for various scale factor A and maximum accelerated energy of protons E p max SED estimated by the best fit model SED of the various assumptions of the spectral index Differential spectrum estimated by the best fit model

9 List of Tables 4.1 Summary of the CANGAROO-II 10-m reflector Summary of the observation periods Terminologies used for the analysis in this thesis Cloud cut and elevation cut conditions and the resulting mean event rate Observation time after the pre-selection Mean scaler counts, cut conditions, and cut ratios in 2002 and Number of excess events showing the gamma-ray signals, statistical significances and event rates assuming the spectral index of the Monte-Carlo simulation is Summary of the the number of excess events with the statistical errors, the acceptances and the differential fluxes with the statistical errors The number of excess events showing the gamma-ray signals, statistical significances and event rates Summary of the the number of excess events with the statistical errors, the acceptances and the differential fluxes with the statistical errors under the assumption of the spectrum with the spectral index of The number of excess events showing the gamma-ray signals, statistical significances and event rates after the conventional cut Ratio of the cut pixels and the acceptance of selecting the bad pixels using ADC distributions. The excess events are counted with α of less than 20. The acceptances are normalized to that without the cut pixels Observation time for Crab nebula before and after the pre-selection Errors of differential fluxes in each region of ADC sum Systematic errors and the uncertainties of the energy in each region of ADC sum Energies and the differential fluxes in each region of ADC sum with all errors and uncertainties Summary of the significances considering the statistical errors and systematic errors

10 Chapter 1 Introduction Cosmic rays were discovered by Hess in Though the studies of cosmic rays have a long history, it is difficult to identify the acceleration sites of cosmic rays since they have electric charges and are diverted by the interstellar magnetic field during the propagation from the source to the Earth. Cosmic rays up to ev are the main component in terms of numbers and are confined in our Galaxy due to the smaller Larmor radius than its disk size. Supernova remnants (SNRs) are believed to be a favored site for accelerating those cosmic rays from the energetics, the energy spectrum, and the chemical composition of the cosmic rays in our Galaxy [7] [34] [43]. The best way to search for the acceleration sites of those cosmic rays is to detect neutral particles, such as gamma-rays and neutrinos, generated from interactions of high-energy cosmic rays with ambient matter, the cosmic microwave background (CMB) radiation or the interstellar magnetic field. Neutrinos, however, are difficult to detect because of their weak interaction. Therefore gamma-rays are the best probe to find the acceleration sites of cosmic rays [62] [42]. If cosmic rays are accelerated up to near ev, gamma-rays in TeV regions, which are not generated in interactions except for those at such high energies, may be emitted. Hence observations of TeV gamma-rays from SNRs are one of the key experiments to explore the origin of cosmic rays. The energy spectra in TeV gamma-rays together with observations at other wavelengths are also the clues to the emission mechanisms. In addition to the Crab nebula, which was the first established source of TeV gamma-rays??, the SNRs which were detected at TeV energies using imaging air Cherenkov telescopes (IACTs) so far are only three; SN1006 [86], Cassiopeia A [3], and RX J [25] [64]. The number of SNRs which emit TeV gamma-rays in our Galaxy can be roughly estimated from energetics to be 100 f from the supernova (SN) rate and the life time of SN, where f is the factor considering the possibility of the detection in TeV gamma-rays and the uncertainties of the estimation. Though the value depends on f, the number of the detections is too small to consider the energetics and establish the hypothesis that the Galactic cosmic rays are mainly accelerated by SNRs. Much more evidences, therefore, are needed. RX J (G ) was one of SNRs which were thought to be detectable in TeV energies with the current IACTs. The hard X-ray spectrum obtained by the ASCA satellite was well fitted by a power-law showing its non-thermal origin [82]. If this emission is the synchrotron radiation from electrons, TeV gamma-rays via the inverse Compton scattering with 2.7K CMB photons may be detected (synchrotron/inverse Compton model). There are not so many SNRs which predominantly emit non-thermal X-rays [55] 8

11 [72] [56] [8]. The radio emission was also found with the Parkes radio-telescope with the power-law spectrum [18] [20]. CO observations showed the richness of large molecular clouds around RX J [58] [61]. They can be targets of proton-nucleon collisions producing gamma-rays. Recently, synchrotron/inverse Compton models considering the different emission regions in X-ray and TeV energies were discussed. When the acceleration and emission regions are different, especially in size, TeV flux could be different from that predicted by the simple synchrotron/inverse Compton model (two-zone model [1]). RX J was appropriate to study the two-zone model because it was the only one that had a large angular size which was estimated by the X-ray observations and had a possibility of the detection of the TeV gamma-rays as described above. From its large angular size, observations of fine structure can be easily carried out. Following these reasons, we selected RX J as one of the most appropriate sources to study the SNR origin of cosmic rays. Hence it was observed by the CANGAROO-II 10-m IACT for two years. The observations, the calibrations, the data analysis and the discussions of the emission mechanisms were carried out by the author. The results of the data analysis and the estimation of the emission mechanisms are reported in this thesis. The origin of cosmic rays, the acceleration theory, the observations of TeV gammarays from SNRs, the details of RX J , the processes of non-thermal emissions are reviewed in Chapter 2. The general methods to detect TeV gamma-rays are described in Chapter 3. The details of CANGAROO-II 10-m telescope are described in Chapter 4. The observations, the calibrations, and the analysis are explained in Chapter 5 and 6. The results of the analysis are presented in Chapter 7. The various checks and the investigations of systematics are described in Chapter 8 and 9. Using these results, we discuss the emission mechanisms and conclude in Chapter 10 and 11. 9

12 Chapter 2 Review 2.1 Origin of Cosmic Rays The origin of cosmic rays has been unclear since Hess discovered the cosmic rays in Before that, the natural radioactivity had been already discovered by Bequerel from the fact that photographic plates became darkened even when fluorescent substances was not exposed to light. The cosmic ray story begins when it was found that electroscopes discharged even if they were kept in the dark well away from sources of natural radioactivity. Hess and Kohlhörster made manned balloon in order to measure the ionization of the atmosphere with increasing altitude. The ionization was 1-2 ion pairs/cc at the see level. The higher they ascended, the more the ionization increased. The ionization at the altitude of 1000m was a few times as much as that at the see level. This was the clear evidence that the radiation came above Earth s atmosphere. This extraterrestrial ionization radiation was called cosmic rays. At the beginning, cosmic rays were thought to be the gamma-rays due to their great penetrating powers. However, it was revealed that the flux of cosmic rays changed with the latitude and they had a tendency to penetrate from the West [4]. From these facts, cosmic rays were found to be the particles with positive electric charges. The balloons with detectors ascended to near the top of the atmosphere in order to detect primary cosmic rays and cosmic rays were revealed to be mainly protons in 1940s. The fluxes of primary cosmic rays were detected using balloon experiments for those below ev and air shower arrays for those above ev. The integral flux is 1/cm 2 /sec/str above 1GeV. Figure 2.1 shows the differential energy spectrum of cosmic rays. The differential spectrum of cosmic rays is well represented by a power-law in the energy range above 1GeV per nucleon. The spectral index is 2.7 for the energies below ev and changes to 3.0 for those above ev (knee). The highest energy of cosmic rays ever detected is ev [44]. Cosmic rays up to ev are the main component in terms of numbers and are confined in our Galaxy due to the smaller Larmor radius than its disk size. The energy distribution is not Maxwellian, i.e. non-thermal. Such particles have extremely large total energy such as 1eV/cm 3 [90]. This is greater than the energy density of the starlight, galactic magnetic field, and cosmic microwave background (CMB), all of which are around 0.3eV/cc. From the point of view of energetics, it is a big problem where and how such a large amount of energy is produced. To investigate the composition of cosmic rays is an alternative way to probe its origin. Roughly speaking, about 99% of the particles are nuclei while about 1% are electrons. Of 10

13 Figure 2.1: Balloon flights of Hess [80]. Preparation for one of his flights in the period (left). Hess after the balloon flights in which the increase in ionization with altitude through atmosphere was discovered (right). Figure 2.2: Differential energy spectrum of cosmic rays [94]. 11

14 those nuclei, about 90% are protons, 9% are α particles, and 1% are the other elements. Figure 2.3 shows the elemental abundances of cosmic rays measured at Earth s orbit compared to the solar system abundances, all relative to silicon [81]. The distribution of Figure 2.3: Cosmic ray elemental abundances measured at Earth compared to the solar system abundances, all relative to silicon [81]: (solid circles) low energy data, MeV/nucleon; (open circles) compilation of high-energy measurements, MeV/nucleon ; (diamonds) solar systems. elemental abundances in cosmic rays is similar to those of typical solar system abundances. Some of the differences give some clues for the origin and propagation mechanism of cosmic rays. The light elements, lithium, beryllium, and boron, are grossly over-abundant in cosmic rays relative to their solar system abundances. These light elements are difficult to produce by the nucleosynthesis both after the big bang and inside the stars. The process of spallations, i.e. high-energy nuclei such as Carbon, Nitrogen, and Oxygen interacting with interstellar matter (mainly protons) during the propagation and producing lighter nuclei, can increase these abundances. For example, following processes produce light nuclei: 12 6 C + p 6 3Li + 4 2He + p + p + n C + p 9 4Be + p + p + p + n +. (2.1) 12

15 Lifetime of cosmic rays is a key to understand the energetics of cosmic rays in our Galaxy. In order to determine lifetime, the ratio between the number of primary particles and those of secondary particles from the above interactions is very useful. Let s consider the spallation of Carbons (C). The ratio between the cross section of producing Boron (B) and the total inelastic cross section is given as σ B = 81.5mb σ total 205mb = 0.4. (2.2) Using the above value, the ratio between the number of Boron and that of Carbon in cosmic rays (B/C ratio) is given as ) C = C p exp ( xλc, B = σ { )} B C p 1 exp ( xλc, σ ( ) total B 1 exp x C = 0.4 λ C ( ), (2.3) exp x λ C where x, C p, and λ C are the column density (g/cm 2 ) that Carbon passed through, the number of the primary Carbon, and the mean free path of Carbon (8.3g/cm 2 ). Figure 2.4 shows the B/C ratio versus kinetic energy [37]. Using B/C is 0.3 from Figure 2.4 Figure 2.4: Spectrum of the ratio between the number of Boron and that of Carbon in cosmic rays as a function of kinetic energy [37]. and Equation (2.4), x 5 g/cm 2 is obtained. Lifetime of cosmic rays is given as T = 5N/c = yr, where N is the Avogadro number, and c is the light speed assuming the matter density is 1 H/cc. On the other hand, radioactive elements in cosmic rays also 13

16 give constraints on the lifetime of cosmic rays. 10 Be is the best element to determine the cosmic-ray life because the lifetime of 10 Be is 10 6 yr. This elements were not positively detected yet. From these considerations, lifetime of cosmic rays τ CR is thought to be 10 7 yr. Using the above lifetime arguments, one may consider the energetics of cosmic rays in our Galaxy. Assuming the region where cosmic rays are confined is the disk with the radius of 10kpc and the thickness of 1kpc, the total energy of cosmic rays in our galaxy is given as 1eV π(10kpc) 2 1kpc ev erg. (2.4) Using Equation (2.4), the required energy for the acceleration of cosmic rays is given as erg 10 7 yr = erg/yr = erg/sec. (2.5) In 1932 Baade and Zwicky had suggested that SNRs were the origin of cosmic rays [7] and Ginzburg and Hayakawa suggested again with more quantitative consideration [34], [43]. Assuming the total shockwave energy of supernova (SN), the rate of SNe, and the efficiency of the acceleration are erg, 0.01 SNe/yr, and 0.1, respectively, the input energy is given as erg/sn 0.01SNe/yr erg/sec. (2.6) It is difficult to give such a large amount of energy except for SNe. From the diffusive shock acceleration theory described in Section 2.3, the acceleration in the shock front of SNR naturally generates the power-law distribution of cosmic rays energy spectrum. The composition of cosmic rays from SNR will be roughly the same as those from the nucleosynthesis inside the stars. From these considerations, SNRs are considered to be the favored sites for accelerating cosmic rays in our Galaxy. 2.2 Fermi Acceleration As discussed in the previous section, SNRs are believed to be a favored site for accelerating cosmic rays up to ev from the energies, the energy spectrum, and the chemical composition of the cosmic rays in our Galaxy. From this section we briefly review the acceleration mechanism. The first idea was introduced by Fermi [27]. Molecular clouds extend to the order of 10 pc with higher density than the interstellar matter. From the Doppler effect of the absorption lines, they move with the dispersion velocity v of 30 km/s. The conductivity in the clouds is so high because their densities are extremely low and also they are highly ionized. The magnetic irregularities are generated from the Alfvén waves due to such a moving plasma in the interstellar magnetic field and generally occupy the interstellar space. Elastic collisions of particles by these magnetic irregularities in the clouds can be considered as if they were elastic collisions against very large mass. Assuming particles collide randomly, the average gain in energy per collision is given as a order of magnitude by (v/c) 2 [27]. The energies of the particles increase statistically. The average particle energy after n times collisions with the clouds is given as E n = E 0 (1 + ξ) n E 0 exp (ξn), (2.7) 14

17 where E 0 is the initial energy of the particle, and ξ is the average energy gain per one collision. The probability that particles make n times collisions is given as P n = (1 P esc ) n, (2.8) where P esc is the probability that particles escape from the accelerated region per collision, which was calculated from the mean free path of the collision with the interstellar matter in Fermi s case. Using Equations (2.7) and (2.8), the number of accelerated particles with the energy of more than E can be calculated as N(E) (1 P esc ) m = (1 P esc) n 1 ( ) δ E, (2.9) P esc where δ is given as m=n δ = ln [1/(1 P esc)] ln(1 + ξ) P esc E 0 P esc ξ. (2.10) This theory leads naturally that the energy spectrum of the accelerated particles obeys an inverse power law. However, ξ is low for this case. This leads to the large power-law index, i.e very soft spectrum considering Equations (2.9) and (2.10). In addition the energy loss may take place at each collision. 2.3 Diffusive shock acceleration (DSA) Instead of Fermi s idea discussed in Section 2.2, more efficient acceleration mechanism by collisions from one direction was introduced, i.e., accelerations in the shock fronts of the SNRs [10], [12]. Figure 2.5 shows a schematic view of particle acceleration around a shock front in the laboratory frame. Interstellar matter in the upstream flows into the downstream through the shock front with velocity v 1 in the rest frame of the shock front, which is greater than the sound speed in the upstream, i.e. faster than the speed for transmitting information, and get slower and denser in the downstream. Suppose that there are the cosmic rays with initial energy of E 1, assuming they are relativistic for simplicity. In the rest frame of the downstream, the energy of the accelerated cosmic ray is given as E 1 = γ v E 1 (1 + β v θ 1 ), (2.11) where a prime ( ) denotes a quantity of the rest frame in the downstream, γ v is the Lorentz factor with velocity v, β is v/c, and θ 1 is the incident angle of the particle, respectively. After multiple elastic scattering with magnetic irregularities, the particle again cross the shock into the upstream with some probability. The energy E 2 is that after the interaction with the downstream medium. The energy gain is given as E E 1 = γ 2 v(1 + β v cos θ 1 β v cos θ 2 β 2 v cos θ 1 cos θ 2) 1, (2.12) where E is E 2 E 1. This value should be averaged over the particle s angle penetrating into the shock front. If the isotropic intensity of the number of particles were given by I, the average of cos θ 1 is given as cos θ 1 = 2π 1 cos θ I cos θd(cos θ) 0 2π 1 I cos θd(cos θ) = 2 3. (2.13) 0 15

18 upstream (outside the SNR) E 1 downstream (inside the SNR) 1 E 2 -v 1 2 -v= -v 1 + v 2 0 shock front x Figure 2.5: Schematic view of the diffusive shock acceleration around the shock front in the laboratory frame. From the same discussion, cos θ 2 = 2/3. From β v 1 of the shock waves in SNRs, the gain of the energy is approximated as E E 1 = 4 3 v 1 v 2. (2.14) c The probability P esc that the scattered particles escape from the accelerated region for each round trip was calculated by Bell [10]. In the rest frame of the shock front the flux of non-thermal particles penetrating into the shock front is given as 1 0 2π d cos θ dφ cρ CR 0 4π cos θ = cρ CR 4. (2.15) The flux of non-thermal particles which escape from the downstream is ρ CR v 2. The P esc is given as P esc = ρ CRv 2 cρ CR /4 = 4v 2 c. (2.16) Using Equation (2.14) and (2.16), the power-law index δ of the integral flux in Equation (2.10) is given as 3 δ = v 1 /v 2 1. (2.17) The compression ratio v 1 /v 2 can be estimated by the dynamics of thermal particles. In the rest frame of the shock front, conservation of mass, momentum, and energy is described as ρ t + (ρv) = 0, (2.18) x 16

19 ρv t + x (ρv2 + P ) = 0, (2.19) { ( )} 1 ρ t 2 v2 + E + { ( ) } 1 ρ x 2 v2 + E v + P v = 0, (2.20) where ρ, v, P, and E are the density, velocity, pressure, and internal energy per unit mass, which is the sum of the kinetic energies of thermal particles, respectively. Assuming a steady state ( / t = 0) and applying these equations to the shock front shown in Figure 2.5, the relations between the physical parameters in the upstream and in the downstream (Rankine-Hugoniot relations) are given as ρ 1 v 1 = ρ 2 v 2 (2.21) ρ 1 v1 2 + P 1 = ρ 2 v2 2 + P 2 (2.22) { ( ) } { ( ) } 1 1 v 1 ρ 1 2 v2 1 + E 1 + P 1 = v 2 ρ 2 2 v2 2 + E 2 + P 2, (2.23) where subscripts 1 and 2 denote the upstream and the downstream, respectively. Equation (2.21), Equation (2.23) reduces to 1 2 v2 1 + E 1 + P 1 ρ 1 = 1 2 v2 2 + E 2 + P 2 ρ 2. (2.24) The plasmas behave as an ideal gas, and using Mayer s relation, E can be written as E = C V T = C V P nrρ = By C V P C P C V ρ = 1 P γ 1 ρ, (2.25) where C V, C P, and γ are the molar heat at constant volume and pressure, and the specific heat, respectively. Using Equation (2.21) and Mach number M v/a = v/ γp/ρ in the adiabatic gas, where a is the sound speed, Equations (2.22) and (2.23) become ( 1 1 r ) γm 2 1 = s 1, (2.26) (1 1r ) M = 2 ( s ) γ 1 r 1, (2.27) where r ρ 2 /ρ 1 = v 1 /v 2 (compression ratio), and s P 2 /P 1, respectively. r and s are given as r = (γ + 1)M 2 1 (γ 1)M From M 1 1 (strong shock) approximation, r change as (2.28) s = 2γM 1 2 (γ 1). (2.29) γ + 1 r = γ + 1 γ 1. (2.30) Assuming γ is 5/3 like monoatomic molecule gas, the compression ratio becomes 4. Using this ratio and Equation (2.17), the index of the integral spectrum of the accelerated particles is unity, i.e. the index of the differential spectrum is 2. This result is consistent with energy spectra of SNRs determined from observations at various wavelength. 17

20 2.4 Observations of TeV gamma-rays from Supernova Remnants Though the researches of cosmic rays have a long history as described in Section 2.1, it is difficult to identify the acceleration sites of cosmic rays since they have electric charges and are diverted by the interstellar magnetic field during their propagation from the source. The best way to search for the acceleration sites of cosmic rays is to detect neutral particles, such as gamma-rays and neutrinos, generated from the interactions of highenergy cosmic rays with the ambient matter, CMB or the interstellar magnetic field. Neutrinos, however, are difficult to detect because of their weak interaction. Therefore gamma-rays are the best probe to find the acceleration sites. This idea was introduced by Morrison and Hayakawa in 1950 s [62], [42]. Below around 10GeV in energy, gamma-rays are totally absorbed by Earth s atmosphere. Hence satellites were launched in order to detect the gamma-rays in such energies. In 1970 s, the SAS-II satellite provided the first detailed information about the gammaray sky and revealed that gamma-ray emission was strongly correlated with galactic disk [40]. Some discrete sources with strong emission such as Crab pulsar, Vela pulsar, and Geminga were also detected. After that, the COS-B satellite was launched and detected more discrete sources [84] as shown in Figure 2.6. Both satellites provided the first detailed Figure 2.6: Number of the sources detected with various detectors versus year. views of the Universe in gamma-rays with energies from about 30MeV to about 5GeV, but the angular resolution was poor, 1, and also the statistics was poor. Therefore it was difficult to identify the gamma-ray emission as known high-energy sources. In 1991, the Compton Gamma Ray Observatory (CGRO) was launched. The Energetic Gamma Ray Experiment Telescope (EGRET) on-board CGRO was 10 to 20 times larger and more sensitive than the previous detectors, with the improved energy range from 20MeV to 30GeV. The angular resolution was strongly dependent on energy but 0.5 at 5GeV was achieved. Five sources including γ Cygni and IC443 were coincident with SNRs [26]. 18

21 Those results might be evidences of the SNR origin of cosmic rays. If so, the energy spectrum of cosmic rays in these SNRs might extend to around knee region. The interactions of cosmic rays at such energies produce gamma-rays of TeV energies. Therefore imaging air Cherenkov telescopes (IACTs) are the most essential detectors to study the origin of cosmic rays, details of which will be described in Chapter 3. However, the observations of six SNRs, including three EGRET sources, by the Whipple Observatory (IACT) gave upper limits on the fluxes above 300 GeV. [13]. The obtained upper limits were below the predicted fluxes based on shock acceleration theory without cutoffs [19] [65]. It turned out that models with cutoffs gave a good fit to the observed spectra [31] [83]. After all the origin of cosmic rays remained unclear. These situations have been drastically changed since intense non-thermal X-ray emission from the rims of Type Ia SNR SN1006 was detected by ASCA as shown in Figure 2.7 [55] [72]. This indicated that electrons were accelerated to energies up to 100 Figure 2.7: X-ray of SN1006 obtained by ASCA satellite [72]. TeV within the shock front. Motivated by the prediction of TeV gamma-rays via inverse Compton emission by these high-energy electrons, observations were carried out with the 3.8m diameter IACT (CANGAROO-I) by the CANGAROO collaboration. Figure 2.8 shows the statistical significance map of gamma-rays from the northeast rim of SN1006 obtained by CANGAROO-I [86]. The gamma-ray peak was coincident with the X-ray maximum. The spectra around TeV energies are different upon the emission mechanism. Figure 2.9 shows a schematic view of the Spectral Energy Distributions (SEDs:E 2 df/de) for synchrotron/inverse Compton model and π 0 decay model. The spectrum of Synchrotron/inverse Compton model is composed of two peaks in SED: synchrotron radiation of electrons at low energies and inverse Compton scattering of electrons with ambient photons at high energies. Decays of π 0 produced by proton-nucleon collisions also produce high-energy gamma-rays. The spectrum of π 0 decay model is trapezoid-shaped in SED. Figure 2.10 shows the SED of SN1006. The SED from radio to TeV can be explained well by the synchrotron/inverse Compton model assuming ambient photons are 2.7K CMB photons. However, from recent theoretical development, π 0 decay model with non-linear acceleration scheme can also explain these spectrum [11]. The main component of cosmic rays is nuclei ( 99%), mainly protons. The de- 19

22 Figure 2.8: Contours of statistical significance map of gamma-rays from SN1006 northeast rim obtained by CANGAROO 3.8m telescope (CANGAROO-I) [86]. Figure 2.9: Schematic view of the Spectral Energy Distributions for synchrotron/inverse Compton model and π 0 decay model. The spectrum of Synchrotron/inverse Compton model is composed of two peaks: synchrotron radiation of electrons at low energies and inverse Compton scattering of electrons with ambient photons at high energies. Decays of π 0 produced by proton-nucleon collisions also produce high-energy gamma-rays. The spectrum of π 0 decay model is trapezoid-shaped. 20

23 10 2 IRAS upper limit EGRET upper limit νsν or ε 2 df/dε [ ev cm -2 s -1 ] radio ROSAT Synchrotron ASCA Inverse Compton B = 4 µg CANGAROO decay parent proton spectrum α=2.2 Emax=5e15 no*e50=2.5 π photon energy ε [ev] Figure 2.10: Spectral Energy Distribution observed from the NE rim of SN1006 [85], where observed fluxes or upper limits of radio [78], infrared, soft X-ray (estimated from Willingale et al. [93]), hard X-ray [72], GeV gamma-rays (calculated from the EGRET archival data), and TeV gamma-rays are presented. Solid lines are the fits based on the model of synchrotron/inverse Compton model and π 0 decay. 21

24 tection of TeV gamma-rays from SN1006 could not be the clear evidence of the proton acceleration. Meanwhile several SNRs were found by the ROSAT all-sky survey. The observations of ASCA revealed intense non-thermal emission from RX J [56], RCW86 [8], and RX J [82] [88]. These SNRs are on the galactic plane. Molecular clouds surrounding them can be a target of accelerated proton interactions. Detections of gamma-rays in sub-tev energies from these SNRs could be the evidences of proton acceleration. The northwest (NW) rim of of RX J was observed by CANGAROO telescope [63] [25]. Figure 2.11 shows the contour map of statistical significance map of RX J northeast rim obtained by CANGAROO 10-m telescope [25]. Figure 2.12 shows the SED from RX J [25]. The spectrum was a good Figure 2.11: Contours of statistical significance map of RX J northeast rim obtained by CANGAROO 10m telescope [25]. match to that from π 0 decay and could not be explained by other mechanisms. Therefore this detection of TeV gamma-rays was thought to be the direct evidence of the proton acceleration. This conclusion, however, was objected by Butt et al. [15] from the point of the density of molecular cloud. Recently the NANTEN observation revealed that a molecular cloud of 200 solar masses was clearly associated with the TeV gamma-ray peak [29] which denied the objection by Butt el al. Reimer and Pohl [77] also claimed the spectrum of the nearest EGRET source was inconsistent with the predicted spectrum by π 0 decay especially in GeV region. The recent non-linear acceleration theory predicted and suggested that the cosmic ray power-law index can be less than 2 and may solve the above problem [11]. Simple synchrotron/inverse Compton model must be considered again if the emission regions in X-ray and TeV energies are different. When the acceleration and emission regions are different, especially in size, TeV flux could be different from that predicted by the simple synchrotron/inverse Compton model (two-zone model [1]). In order to explain the TeV emission of RX J with this scheme, the volume ratio of V T ev /V X = 1000 is necessary [73]. There is, however, no other physical evidence for this value. The HEGRA group detected TeV gamma-rays from Cassiopeia A using a stereoscopic Cherenkov telescope system [3]. From the hard X-ray continuum, a lower limit to the average magnetic field was estimated to be 0.5 mg [89]. Therefore the TeV spectrum is difficult to explain by inverse Compton scattering of electrons due to such a high magnetic field. By this discussion Cassiopeia A is thought to be the proton acceleration site. However the evidences for proton acceleration are still sparse and are 22

25 E 2 df/de or E F(>E) (ev cm -2 s -1 ) radio ASCA EGRET CANGAROO 10-5 Photon energy, E (ev) Figure 2.12: Spectral Energy Distribution from RX J , and emission models [25]. The radio data was obtained by ATCA [21]. The shaded area between thick lines shows the ASCA GIS data. The EGRET upper limit corresponds to the flux of 3EG J [41]. The TeV gamma-ray points show CANGAROO data. Lines show model calculations: synchrotron emission (solid line), inverse Compton emission (dotted lines). bremsstrahlung (dashed lines) and emission from π 0 decay. Inverse Compton scattering and bremsstrahlung are plotted for two cases: 3µG (upper curves) and 10µG (lower curves). 23

26 not conclusive. In order to establish the hypothesis that the Galactic cosmic rays are mainly accelerated by the SNRs, much more evidences are needed. Other SNRs with non-thermal X-ray emission should be observed with IACTs, such as RCW86 and RX J , which have been already observed by CANGAROO. The data analysis and results of RX J were reported in this thesis. The details of RX J observations were described in the next section. Further studies will be also carried out by the next generation Cherenkov telescopes with lower energy threshold such as CANGAROO-III. 2.5 SNR RX J (G ) RX J (G ) is a SNR located at the southeast corner of the Vela SNR. It was discovered at X-ray energies during the ROSAT all-sky survey by Aschenbach [5] shown in Figure Its apparent size was around 2. The MeV 44 Ti line was Figure 2.13: Soft X-ray images of Vela SNR observed by ROSAT. This is for photon energies < 1.3 kev. RX J is on the lower left. detected with COMPTEL by Iyudin et al. [49]. 44 Ti decays into 44 Sc emitting two hard X-ray lines of 68 kev and 78keV. The lifetime of 44 Ti is 60yr [67] [36]. From these values, the weighted mean of the lifetime was derived to be 90.4±1.3 years [49]. The effective 44 Ti lifetime could be larger, depending on the degree of ionization of the 44 Ti and its Lorentz factor. 44 Sc decays further to 44 Ca while emitting a gamma-ray line at 1.156MeV with the lifetime of 3.9 hr. By combining the gamma-ray line flux and the X-ray diameter with an assumed typical 44 Ti yield, and taking as representative an expansion velocity of 5000km s 1 for the supernova ejecta, the distance and age were estimated to be 200 pc and 680yr, respectively [49]. It was not recorded historically. This may have been seen in measurements of nitrate abundances in Antarctic ice cores [14]. Supernovae 24

27 can produce NO 3 when their radiations ionize the molecules in the atmosphere. The X- ray emission line at 4.1±0.2 kev was only detected in the northwest shell by ASCA [88]. This line was thought to come from highly ionized Ca. We, however, cannot distinguish among the Ca isotopes using X-ray data. Assuming that most of the Ca is 44 Ca, the age of RX J was estimated to be around 1000 yr combining the amount of 44 Ca and the observed flux of the 44 Ti [88]. Aschenbach, Iyudin, and Schönfelder estimated the distance and age again [6]. They estimated the expansion velocity using X-ray spectra at the limb obtained by ROSAT. The minimal, best-estimate, and maximal expansion velocities were 2000, 5000, and km s 1, respectively. Model calculations provide a range for the mass yield of 44 Ti. Considering these uncertainties, the upper limit of the distance of RX J was estimated to be 500pc and 1100yrs for the age. Chen and Gehrels have also used the X-ray temperature obtained from ROSAT data for the central region to derive a range of km s 1 [17]. If this is true, the remnant is currently expanding too slowly to be caused by a Type Ia supernova. The estimation of Aschenbach, Iyudin, and Schönfelder, however, allow it to be a Type Ia for the expansion speed. The central region of the SNR was observed with ROSAT [5], ASCA [82], BeppoSAX [59], and Chandra [74]. Based on the X-ray-to-optical flux ratio, the X-ray source in the central region was likely the compact remnant of the supernova explosion that created the RX J Figure 2.14 shows the hard X-ray images of RX J observed by ASCA GIS [82] [88]. The images clearly shows shell- Figure 2.14: Hard X-ray image of RX J observed by ASCA GIS (E = kev) [82]. The image consists of a mosaic of seven individual fields. Contours represent the outline of the Vela SNR as seen in ROSAT survey data with the PSPC. like morphology. The hard X-ray spectrum was well fitted by a power law. The matter density of the X-ray peak was estimated to d 1/2 1 f 1/2 [H/cm 3 ], where d and f are the distance normalized to 1kpc and the filling factor, respectively [82]. While simple 25

28 scaling of the column density to estimate the distance was clearly rather uncertain, it appears that the remnant is at least several times more distant than Vela. The distance to the Vela SNR was estimated to be 250±30 pc using Ca II and Na I absorption line spectra toward the OB stars in the direction of Vela SNR [16]. The distances to the OB stars were well determined using trigonometric parallaxes and spectroscopic parallaxes based on photometric colors and spectral types. The radio emission was found with the Parkes radio-telescope [18] [20]. The fluxes at 2.42 and 1.40 GHz were 33±6Jy and 40±10 Jy, respectively [20]. The spectral index were 0.40±0.15 at the northern section of the shell [20] Figure 2.15 shows the 4.85 GHz radio images [20]. Shell-like morphology can Figure 2.15: Radio images at 4.85GHz observed by the Parkes radio-telescope, centered on RX J The angular resolution is 5, and the rms noise is approximately 8 mjy beam 1. The grey-scale wedge is labelled in units of Jy beam 1. The black circle is centered on the X-ray coordinates of the source and is 1.8 in angular diameter. be seen. But the confusing structures from Vela SNR exist. CO observations showed the richness of large molecular clouds around RX J in the Vela Molecular Ridge [58]. They can be targets of proton-nucleon collisions. The detailed morphology was mapped with the NANTEN 4 m milli-metre radio telescope [61]. Figure 2.16 shows the CO map around the Vela SNR. CO observations has a better accuracy to determine the distance than 21cm radio observations because of their narrow Doppler broadening. The correlation between RX J and the molecular clouds were not yet investigated. From above observations, the characteristics of RX J are similar to those of RX J TeV electrons can be expected from the non-thermal X-ray emission. If the ambient magnetic field is not so strong, TeV gamma-rays from inverse Compton scattering are produced. On the other hand, TeV gamma-rays from π 0 decay produced by proton-nucleon collisions may be detected because the nearby molecular clouds exists. 26

29 Figure 2.16: Integrated intensity map of CO obtained with NANTEN 4 m milli-metre radio telescope (contour) and a soft X-ray image by ROSAT (gray scale) [61]. The cross indicates the position of the Vela pulsar. The equatorial coordinates are indicated by the dashed lines. 2.6 Processes of Non-thermal Emissions The processes of non-thermal emissions in SNRs were briefly reviewed in this section. More detailed calculations of the gamma-ray spectrum are described in Chapter π 0 decay π 0 s are produced in collisions of accelerated nuclei which are mainly protons, with interstellar matter which also consists of nuclei, mainly protons. The π 0 s immediately decay in two gamma-rays within the mean lifetime of γ π seconds, where γ π is the Lorentz factor of the π 0 s. These gamma-rays have similar energy spectrum to that of the parent high-energy particles because of the scaling hypothesis. Details of calculations of differential cross section is given in Appendix B Synchrotron Radiation Photons are emitted from charged particles accelerated by the Lorentz force in the magnetic field. When charged particles are relativistic, the frequency spectrum can extend to many times the gyration frequency. This radiation is known as synchrotron radiation. The following is the simple estimation of the total emitted power of the electron [68] [79]. The Lorentz force is produced by only v which is the velocity of the electron perpendicular to the direction of the magnetic field. Therefore we should see only v. We assume v is zero, where v is the velocity of the electron parallel to the magnetic field. The magnetic field B can be regarded as the electric field E in the rest frame of the electron. This can be derived using Lorentz transformation as E = γβ B, (2.31) 27

30 where γ and β are the Lorentz factor of the electron and v /c (c is light speed), respectively. Using the Larmor s formula, the total emitted power in the rest frame is given as ( P = 2e2 e γβ ) 2 B 3c 3 m B 2 = 2σ T cγ 2 β 2 8π, (2.32) where e and m are the charge of the electron and the rest mass of the electron, respectively. For an isotropic distribution of velocities it it necessary to average this formula over all angles for an given speed β. Let α be the pitch angle, which is the angle between the magnetic field and the velocity. Then we obtain β 2 = β 2 sin α 2 dω where Ω is the solid angle, respectively. And the result is 4π = 2β2 3, (2.33) P = 4 3 σ Tcγ 2 β 2 U B, (2.34) where σ T and U B are the cross section of Thomson scattering and the energy density of the magnetic field, respectively. The total emitted power P is the Lorentz invariance and is preserved under Lorentz transformation. Hence P is give as P = 4 3 σ Tcγ 2 β 2 U B. (2.35) Synchrotron radiation is important only for electrons since P is proportional to 1/m 2 for high-energy particles from Equation (2.35). The frequency spectrum can extend to many times the gyration frequency. Figure 2.17 shows the spectral distribution of the power of the total (over the directions) radiation from charged particles moving in a magnetic field as a function of ν/ν C [35], where ν and ν C are the frequency of the emitted photons and ν C = 3eBγ 2 /4πmc. The spectrum has a roughly monochromatic peak Inverse Compton Scattering When relativistic electrons move in the photon field, the scattered photons by Compton scattering gains its energy from the electron. This process is called inverse Compton (IC) scattering. The energy of the ambient photon in the rest frame of the electron is given as hν = γhν(1 + β cos θ), (2.36) where h, ν, γ, ν, β, and θ are Planck constant, the frequency of the photon in the laboratory frame, the Lorentz factor of the electron, the frequency of the electron in the rest frame of the electron, the velocity of the electron, and the angle of incidence from the direction of the electron motion in the laboratory frame, respectively. Assuming the distribution of the ambient photons is isotropic, the photon energy is γhν. In case of γhν mc 2, where m is the rest mass of the electron, the scattering in the rest frame of 28

31 Figure 2.17: Spectral distribution of the power of the total (over the directions) radiation from synchrotron radiation [35]. ν and ν C are the frequency of the emitted photons and ν C = 3eBγ 2 /4πmc. the electron is approximately Thomson scattering, i.e. elastic. Hence the energy of the scattered photon is given as hν = γhν (1 + cos ϕ) γ 2 hν, (2.37) where ν and ϕ are the frequency of the scattered photon in the laboratory frame and the scattering angle of the photon in the rest frame, respectively. The energy of the scattered photon exactly averaging over angles is given as h ν = 4 3 γ2 hν. (2.38) Using Equation (2.38), the energy loss rate of the electron is given as de dt = 4 3 σ Tcγ 2 U, (2.39) where σ T, and U are the cross section of Thomson scattering and the energy density of the radiation field nhν (n is the number density of ambient photons), respectively Bremsstrahlung When charged particles are passed through the Coulomb field of a nucleus, photons are emitted. This is called bremsstrahlung. As described in Subsection in detail, the cross section and the emitted power of bremsstrahlung are proportional to 1/m 2, where m is the rest mass of the charged particle. Therefore bremsstrahlung is important for electrons and is negligible for nuclei. 29

32 Chapter 3 Imaging Air Cherenkov Technique 3.1 Overview As was discussed in the previous chapter, the detection of TeV gamma-rays is a good way to study particle accelerations. However, it is not easy to detect significant signals in TeV energies due to the small statistics at such energies. Besides the depth to stop all particles produced in the cascade above TeV energies is too large for the satellite. Hence TeV gamma-rays penetrate into the atmosphere without being detected in the space. The earth, however, has sufficient material, i.e. the atmosphere itself. Therefore very large-scale detector can be made using the atmosphere as a part of detector. When the gamma-rays and cosmic rays penetrate into the atmosphere, they interact with molecules of air and generate electromagnetic and/or hadronic cascade, respectively. This phenomenon is well known as Extensive Air Showers (EASs). EASs develop and the number of the particles in EASs reaches the maximum (shower max) approximately at an order of 10km from the sea level. It is impossible to detect such huge EASs by calorimetric ways like satellites because of low flux. Instead of this method, the optical lights from the showers, such as fluorescence and/or Cherenkov radiation can be used as information of showers. Fluorescence is isotropic radiation. Hence it is difficult to detect it on the ground due to its weakness for TeV gamma-rays. On the other hand, Cherenkov radiation becomes strongly peaked in the direction of particle motion, i.e. roughly along the axis of EASs. It spreads on the ground with an order of 100m. If the detectors have the large area collecting photons such as 10m 2, it is possible to detect such a light as a short-time pulse of 10nsec. Furthermore, the extremely larger effective area ( m 2 ) than the satellite can overcome small statistics. It, however, is still difficult to detect TeV gamma-rays without removing showers which are generated from cosmic rays. Fortunately developments detecting profiles of showers can distinguish gamma-rays from cosmic rays. Imaging air Cherenkov technique can detect the difference in the images of the photons on the ground and distinguish them. CANGAROO telescope is one of the imaging air Cherenkov telescope (IACT) for realizing such a technique. The details of the EASs, Cherenkov radiation, and imaging air Cherenkov technique are described in the following sections. 30

33 3.2 Extensive Air Showers When gamma-rays and cosmic rays in TeV energies penetrate into the atmosphere, they interact with the atmosphere. Showers are classified by the primary particles, i.e electromagnetic particles (electrons, positrons, and gamma-rays) and hadrons (mainly protons and nuclei) Electromagnetic Showers When the energetic electrons pass through the matter, they emit photons due to their acceleration in the Coulomb field of the atomic charge of the nuclei (bremsstrahlung). Each of the secondary photons then reproduces electron-positron pairs through the pair creation process. As a result of these interactions, the number of both electrons and photons increases. This phenomenon is called electromagnetic shower. The probability of the photon emission and that of pair creation were calculated by Bethe and Heitler [45]. The probability for emission of a photon in the energy interval (E, E + de ) by an electron of energy E after traversing a medium of thickness dx g/cm 2 is given as Φ de E = 4 N 137 A z2 r0 2 de E E 4 = N 137 A z2 r0 2 dv v [{ 1 + ( ) 2 ( ) } ( ) ] 1 E E E E log (191Z 1 3 ) + 1 E 9 1 E E [ {1 ] + (1 v) 2 2 (1 v)} log (191Z 1 3 ) + 1 (1 v), (3.1) 3 9 where Z and A are the atomic number, and atomic weight of the traversed matter, respectively, N is Avogadro s number, r 0 is the classical electron radius e 2 /mc 2, and v E /E. Figure 3.1 shows the spectrum of bremsstrahlung approximately calculated from Equation (3.1). From Equation (3.1), the energy loss of bremsstrahlung is approximated 140 f(x) Arbitrary unit v=e /E Figure 3.1: Spectrum of bremsstrahlung approximately calculated from Equation (3.1), where E, E, and v are the electron energy, emitted photon energy, and E /E, respectively. The probability is 1/v at v 1. as where we put de E dx = E Φ de 0 E X 1 E, (3.2) X 1 = 4 N 137 A z2 r0 2 log Z 1 3. (3.3) 31

34 X is called radiation length, which is 37 g/cm 2 in the air. The probability of pair creation by a photon of energy E generating an electron in (E, E + de ) is again approximated as Ψ de E = 4 N 137 A z2 r0 2 de E = N A z2 r 2 0dv [{ ( E ) 2 ( + 1 E E E ) E 3 E ( ) } ( ) ] 1 E log (191Z 1 3 ) 1 E 9 1 E E [ {(v) 2 + (1 v) v (1 v)} log (191Z 1 3 ) 1 9 v (1 v) ]. (3.4) Figure 3.2 shows the spectrum of pair creations approximately calculated from Equation (3.4). Using Equation (3.4), the total probability of pair creation is given as E E 1 f(x) Arbitrary unit v=e /E Figure 3.2: Spectrum of pair creations approximately calculated from Equation (3.4), where E, E, and v are the photon energy, created electron energy, and E /E, respectively. E 0 Ψ de E 7 9 X 1. (3.5) Except for the above two processes, the ionization loss and the multiple Coulomb scattering effect must be considered. Ionization loss can be neglected if the electron has a energy of more than the critical energy ( 700/Z MeV; Z is the atomic number). In case of the air, the critical energy is 81 MeV. Assuming the ionization loss and multiple Coulomb scatterings are neglected, the pair creation and the Bremsstrahlung occur once per radiation length. When the electron energy becomes near the critical energy, the ionization loss becomes dominant and the shower development stops. In order to estimate the shower development, we adopt a simple model where the energy dissipation of the particles occurs only through the constant ionization loss. The equation is given as T ɛ 2 t dt = E 0, (3.6) 0 where T, ɛ and E 0 are the depth at the shower maximum, the ionization loss per radiation length, and the energy of primary particle, respectively. Here we assume that the number of particles in showers is 2 t at depth t. From Equation (3.6), T and E 0 are given as ( ) E0 T ln, (3.7) ɛ N max 2 T E 0, (3.8) 32

35 where N max is the number of particles in the shower maximum. More practical forms are given as ( ) E0 T ln MeV 2, (3.9) ( ) E0 N max (3.10) 1T ev T and N max are 10 and 1000 at 1TeV, respectively. At these processes in the showers, the transverse momenta of the secondary particles are considered to be an order of m e c/2 (i.e. relativistic beaming). For example, the Lorentz factor of the electron with the critical energy is 160 and the emission angle is 1/γ radian As was described above, particle interactions in the matter are characterized by X (g/cm 2 ), which is the depth from the top of the atmosphere. Figure 3.3 shows the definitions of the variables representing the height and depth of the atmosphere. X V is defined X=0 Top of the atmosphere X V X S h V h S Shower max (X ~10) V h=0 On the ground Figure 3.3: Definitions of the variables representing the height and depth of the atmosphere. X is the depth from the top of the atmosphere (g/cm 2 ). h is the distance from the see level (cm). The shower max of gamma-rays at around 1 TeV is X V 10. as X V (h V ) = h V ρ(h )dh, (3.11) where ρ(h) is the density of the atmosphere at the height of h. Assuming the atmosphere is the ideal gas and using the pressure p(h V ) = X V (h V ) and the density ρ(h V ) = dx V /dh V, the equation of state is written as p ρ = X V dx V /dh V = RT, (3.12) where R is Rydberg constant. From Equation (3.12), the depth X V is given as X V = X 0 exp ( h V /h 0 ), (3.13) 33

36 where X g/cm 2, and h 0 = RT is the scale height. The depth of the atmosphere is roughly parameterized by Equation (3.13) and the scale height h 0 (T ). For example, X V s are 1X, 7X and 10X at 23km, 10km and 8km from the sea level, respectively. Figure 3.4 shows a schematic view of electromagnetic showers of 1TeV gamma-rays in the air. In summary, gamma-rays in 1TeV energies penetrate into the air. First Elevation (km) Depth 1X 7X 10X n= Depth e + e - 1X 2X 3X photon 0 50m 28X n= Figure 3.4: Schematic view of electromagnetic showers of 1TeV gamma-rays in the air. X and n are the radiation length of the electron in the air and the refractive index. interaction occurs at a point of 20km from the see level, corresponding to one radiation length. The showers do not develop quickly because of the low pressure there. The shower size suddenly increases exponentially around altitude of 10 km. Their whole shapes are thin. After the shower maximum, particles lose their energy due to the ionization loss and stop shower development Hadronic Showers The mean free path of nucleons in air in TeV energies is 100g/cm 2 (inelastic collision length). The first inelastic collision occurs at a height of 16 km according to the U.S. Standard Atmosphere table. The first proton-proton collision produces pions and secondary nucleons. Figure 3.5 shows a schematic interaction processes of hadronic showers [92]. Typical transverse momentum of the secondary particle is MeV, which is an order of magnitude larger than the case of electromagnetic showers. Therefore the 34

37 Figure 3.5: Schematic interaction processes of hadronic showers [92]. 35

38 opening angle of the shower development is larger than that of the electromagnetic showers. Experimentally shower max is known to be located around 3-inelastic interaction length, i.e., 10 km altitude 300 g/cm 2. The secondary π 0 s have very short life times, sec, before decaying to two gamma-rays. The secondary nucleons and charged pions proceed to the next collisions with nucleons in the air, which produce pions and secondary nucleons again until their energies drop below those required for multiple pion production, i.e. about 1GeV. Below 1GeV, the secondary protons around 1GeV lose their energy due to the ionization loss and its decay. The charged pions decay to muons and muon neutrinos via π + µ + + ν µ, π µ + ν µ. (3.14) The life times of charged pions are 10 8 sec. The low energy muons decay after 2µsec to positrons, electrons and muon neutrinos via µ + e + + ν e + ν µ, µ e + ν e + ν µ. (3.15) The produced positrons and electrons form the electromagnetic showers. The high-energy muons ( 2GeV) reach the Earth s surface because its interaction is weak. The hadronic showers have the extended structure because of the above mentioned reasons, compared to the electromagnetic showers. Figure 3.6 shows the showers of gamma-rays and cosmic rays in 1TeV simulated by the Monte Carlo methods. Figure 3.6: Showers of gamma-rays (left) and protons (right) in 1TeV simulated by the Monte Carlo methods. 36

39 3.3 Cherenkov Radiation Blackett in 1948 predicted that the Cherenkov light from EASs should be detectable from the surface on the earth. It was confirmed observationally several years later by Galbraith and Jelley [32]. We introduce the characteristics of the Cherenkov radiation in this section. Suppose an electron is moving faster than the light in the medium which is a perfect isotropic dielectric material. Figure 3.7 shows a schematic view of Cherenkov radiation. Using Huygens s principle, the wave front is determined as shown in the Figure c t + - E + - v t - C + c =c/n e - E=0 Figure 3.7: Schematic view of Cherenkov radiation, where n, c, c, v, t and E are the refractive index, the light speed in the vacuum, the light speed in the dielectric medium, the velocity of the electron, a time and electric field, respectively. 3.7 and cos θ C is given as cos θ C = 1 βn, (3.16) where v, β and n are the velocity of the electron, v/c and the refractive index, respectively. At the back of the wave front, the medium is polarized by the electric field. The atoms in the medium behave like dipoles. Each dipole radiates a short electromagnetic pulse. From Equation (3.16), θ C of the electrons of the critical energy in the air is 0.7 degree at 10km from the sea level, where n is On the ground, the distributions of the above photons are the circle with the radius of 10km tan (0.7 ) 120m. This characteristic makes the effective area extremely large ( m 2 ). Figure 3.8 shows the examples of the images and the lateral distributions of photons produced from gamma-ray showers in 1TeV. Figure 3.9 shows the examples of protons in 3TeV. Frank and Tamm estimated the light yield of Cherenkov radiation by the classical theory [28]. The radiation energy by an electron of angular frequency ω after traversing a medium of thickness dl g/cm 2, is 37

40 Figure 3.8: Images (left) and the lateral distributions (right) of photons produced from gamma-ray showers in 1TeV. Figure 3.9: Images (left) and the lateral distributions (right) of photons produced from proton showers in 3TeV. 38

41 given as dw dl = e2 (1 1 )ωdω, (3.17) c 2 βn>1 β 2 n2 Cherenkov radiation is independent of ω in terms of where e is the electron charge. the number of photons because it does not have any specific frequency. Therefore the number of photons emitted from Cherenkov radiation is proportional to dλ/λ 2, where λ is the wavelength of the emitted photons. The number of photons emitted by an electron between wavelengths λ 1 and λ 2 is given from Equation (3.17) as N = 2παl( 1 1 )(1 1 ), (3.18) λ 1 λ 2 β 2 n2 where α and n are the fine structure constant = e 2 / c = 1/137 and the average refractive index of the medium. From Equation (3.18), the number of photons emitted from the relativistic electron in the air per meter is 10 at the wavelengths between 300nm and 600nm corresponding to the range typically covered by the photomultiplier tubes (PMTs). The atmosphere, however, scatters and absorbs these photons. Figure 3.10 shows the transmission of air. Including all effects, the photon spectrum peaks at around 300nm. Figure 3.10: Transmission of the Cherenkov photons. The detectors of IACTs must cover such a energy range. 3.4 Imaging Air Cherenkov Technique Most of the showers are generated from cosmic rays. In order to distinguish between hadronic showers and electromagnetic showers, the characteristics of directional distribu- 39

42 tions of photons at the ground provides useful information. The photon distributions are different due to the shower developments as were described in the Section 3.2. Figure 3.11 shows examples of the distributions of photons on the focal plane of IACTs. Hillas Figure 3.11: Examples of the distributions of photons on the camera plane of IACTs. introduced the parameters which characterizes the shapes of the photon images (Hillas parameters) [46]. The most important parameter α (image orientation angle) was introduced by Punch et al [75]. Figure 3.12 shows a schematic view of Hillas parameters and α. The root mean square (RMSs) spreads of light along the major axis and the minor Length Width Alpha Distance Source Centroid Figure 3.12: Hillas parameters [46] and α [75]. axis are Length and Width, respectively. Distance is the distance between the centroid of the image and the source position. α is the angle between the major axis and the line which connects the source position to the centroid of the image. Asymmetry is the cubic root of the third moment of the image along the major axis. The details are described in Appendix A. As shown in Figure 3.11, the gamma-ray images are compact compared with proton images. Therefore the selection of the compact images using Length and Width is effective. 40

43 The most effective way, however, is the so-called α cut. The gamma-ray images should be concentrated near α = 0 because the shower axis is parallel to the pointing direction of the telescope. On the other hand, the axes of the showers generated from cosmic rays point to various directions and the images are not concentrated to the specific value of α. Figure 3.13 shows the schematic view of images generated from gamma-rays and protons. Whipple group first detected Crab nebula using above parameters [76] and verified the Figure 3.13: Schematic view of images generated from gamma-rays and protons. capability of IACTs. Figure 3.14 shows the example of the α distributions obtained by the Whipple Observatory [48]. The CANGAROO-II telescope has a 10-m reflector as large as the Whipple telescope described in the next chapter. Most important difference between these two is that Whipple is located in the northern hemisphere and we are in the southern hemisphere. 41

44 Figure 3.14: Distributions of α about Markarian 421 obtained by the Whipple Observatory [53]. The solid line and the dotted line show the ON-source data and the OFF-source data, respectively. 42

45 Chapter 4 The CANGAROO-II 10-m Telescope The CANGAROO (Collaboration of Australia and Nippon (Japan) for a GAmma Ray Observatory in the Outback) 1 is the international collaborated experiment for very highenergy gamma-ray observations using IACTs since The observation site is located near Woomera, South Australia ( E, S, 220 m a.s.l.). The CANGAROO-II project is exploring the southern sky at gamma-ray energies of TeV. Figure 4.1 is a picture of the CANGAROO-II 10-m telescope. Before proceeding to this project, we used CANGAROO-I telescope with the 3.8m diameter reflector [39] and detected TeV gamma-ray emission such objects as pulsar nebulae (PSR [54], the Crab [87]), and SNRs (SN1006 [86], RX J [64]). The CANGAROO-II 10-m telescope had been in operation since April, 2000, and detected SNR RX J [25] and the active galactic nuclei Mrk 421 [71], and the starburst 1 The collaborators of CANGAROO group: A. Asahara 1, G.V. Bicknell 2, R.W. Clay 3, Y. Doi 4, P.G. Edwards 5, R. Enomoto 6, S. Gunji 4, S. Hara 6, T. Hara 7, T. Hattori 8, Sei. Hayashi 9, C. Itoh 10, S. Kabuki 6, F. Kajino 9, H. Katagiri 6, A. Kawachi 6, T. Kifune 11, L.T. Ksenofontov 6, H. Kubo 1, T. Kurihara 8, R. Kurosaka 6, J. Kushida 8 Y. Matsubara 12, Y. Miyashita 8, Y. Mizumoto 13 M. Mori 6, H. Moro 8, H. Muraishi 14 Y. Muraki 12 T. Naito 7 T. Nakase 8, D. Nishida 1, K. Nishijima 8, M. Ohishi 6, K. Okumura 6, J.R. Patterson 3, R.J. Protheroe 3, N. Sakamoto 4, K. Sakurazawa 15, D.L. Swaby 3, T. Tanimori 1, H. Tanimura 1, G. Thornton 3, F. Tokanai 4, K. Tsuchiya 6, T. Uchida 6, S. Watanabe 1, T. Yamaoka 9, S. Yanagita 16, T. Yoshida 16, T. Yoshikoshi 17 (1) Department of Physics, Graduate School of Science, Kyoto University, Sakyo-ku, Kyoto , Japan (2) Research School of Astronomy and Astrophysics, Australian National University, ACT 2611, Australia (3) Department of Physics and Mathematical Physics, University of Adelaide, SA 5005, Australia (4) Department of Physics, Yamagata University, Yamagata, Yamagata , Japan (5) Institute of Space and Astronautical Science, Sagamihara, Kanagawa , Japan (6) Institute for Cosmic Ray Research, University of Tokyo, Kashiwa, Chiba , Japan (7) Faculty of Management Information, Yamanashi Gakuin University, Kofu, Yamanashi , Japan (8) Department of Physics, Tokai University, Hiratsuka, Kanagawa , Japan (9) Department of Physics, Konan University, Kobe, Hyogo , Japan (10) Ibaraki Prefectural University of Health Sciences, Ami, Ibaraki , Japan (11) Faculty of Engineering, Shinshu University, Nagano, Nagano , Japan (12) Solar-Terrestrial Environment Laboratory, Nagoya University, Nagoya, Aichi , Japan (13) National Astronomical Observatory of Japan, Mitaka, Tokyo , Japan (14) School of Allied Health Sciences, Kitasato University, Sagamihara, Kanagawa , Japan (15) Department of Physics, Tokyo Institute of Technology, Meguro, Tokyo , Japan (16) Faculty of Science, Ibaraki University, Mito, Ibaraki , Japan (17) Department of Physics, Osaka City University, Osaka, Osaka , Japan 43

46 Figure 4.1: CANGAROO-II 10-m telescope. 44

47 galaxy NGC253 [47]. The details of the CANGAROO-II 10-m telescope is described in the following sections. The observations of CANGAROO-III stereoscopic system with two telescope started in December, CANGAROO-II 10-m telescope is defined as the first one of the CANGAROO-III telescopes. The full four telescopes will be in operation in Reflector IACTs have a reflector to collect Cherenkov photons from EASs. The number of these photons is proportional to the energy of the incident gamma-ray. In order to obtain the lower threshold, the large reflector should be equipped with the IACT. Its size is limited by the engineering and the cost. Such a large effective area can be obtained by multiple segmented mirrors. Table 4.1 shows the summary of the CANGAROO-II 10-m reflector [52]. The mount is alt-azimuth which is typically adopted to IACTs. The diameter is 10.4m Frame parabolic Diameter 10.4m Focal length 8m f 0.77 Number of segmented mirrors 114 Mirror diameter 80cm Total collecting area 57.3 m 2 Mirror segment shape spherical Mirror curvature 16.4m Mirror material CFRP Table 4.1: Summary of the CANGAROO-II 10-m reflector. as large as the other major reflector of the IACT such as Whipple. Most of the IACTs adopt Davis-Cotton type of reflector, which have a frame of spherical shape with multiple spherical segmented mirrors. Davis-Cotton type has a better off-axis focusing property, but the timing information is not good (the maximum time variation of photon 6nsec at (f/0.7)) due to the optical path differences. On the other hand, the CANGAROO-II 10-m telescope adopt a parabolic reflector. The arrival time information can be used (the maximum time variation of photon 0.2nsec at (f/0.7)) and the Night Sky Background photons can be rejected using the narrow timing gate described in Chapter 6. If the mirror curvature is chosen appropriately, an acceptable off-axis performance can be achieved. In the case of CANGAROO-II 10-m telescope, a 16.4-m radius of curvature gives a best performance. The above mentioned segmented mirrors are made of Carbon Fiber Reinforced Plastic (CFRP) with aluminum sheet in order to reduce weight. The diameter, the thickness and the weight are 0.8m, 18mm and 5.5 kg, respectively. The average density is about one fifth of the ordinary glass mirror ( g.cm 3 ). A schematic cross section of a segmented mirror is shown in Figure 4.2. The combination of low density, high shearstrength foam and the prepregs (sheets of carbon fiber impregnated with resin) achieves the light segmented mirror which is hard to deform [52]. The deflections by the gravity 45

48 Figure 4.2: Schematic illustration of the cross section of a segmented mirror. The prepreg is made of carbon fibers impregnated with resin. were estimated to be as small as a few µm at the edge of the mirror. In reality, the curvature radius varies from 15.9 m to 17.1m. The small curvature mirrors were placed in the inner region and the large curvature ones in the outer region in order to minimize the aberration. Point spread function of each segmented mirror is 0.1 (FWHM). This value is included in Monte Carlo simulation as was described in Subsection The reflectivity is dependent on the wavelength. It is 80% between 300 and 600nm where Cherenkov photons after the transmission in the air are mainly distributed. 4.2 Imaging Camera The camera of IACTs is equipped on the focal plane and detects the photons from the showers. The imaging camera of CANGAROO-II 10-m telescope is shown in Figure 4.3. The camera consists of 552 PMTs which cover field of view. Figure degree Figure 4.3: Imaging camera of CANGAROO-II 10-m telescope. shows the PMT (Hamamatsu R4124UV) used for pixel. The PMT has a 13mm (1/2 46

49 Figure 4.4: PMT (Hamamatsu R4124UV) used for the CANGAROO-II 10-m telescope. inch) diameter with UV glass window. The photocathod is made of bialkali. Its spectral response is shown in Figure 4.5. The quantum efficiency is 20% between 300nm and Figure 4.5: Spectral response of photocathod. 500nm where Cherenkov photons after the transmission in the air are mainly distributed. Each PMT covers But the areas of each PMT are not entirely sensitive. Figure 4.6 shows the light guides of CANGAROO-II 10-m telescope. These are attached in front of the PMTs in order to cover the dead space of the PMTs. A group of 4 PMTs (1 4 array) are installed in one amplifier card (LeCroy TRA402S). 4 cards are installed in one high voltage supply box (LeCroy 1461N). Each group of 16 PMTs is supplied with the same voltage of about 700V, and their gains are adjusted within 15%. These units of 16 PMTs are arrayed in a 6 6 square. Each corner module consists of only 10 PMTs. 47

50 Figure 4.6: Light guides of CANGAROO-II 10-m telescope. 4.3 Electronics and Data Acquisition System A block diagram of data acquisition (DAQ) for the CANGAROO-II 10-m telescope is shown in Figure 4.7. The signals from all PMTs are fed to the electronics hat by 36m twisted pair cables. The cables are connected to a VME divider module and a HVpower supply. The VME divider module divides the signal from PMTs to both the VME 9U-bus 32ch 12bit charge ADC (HOSHIN 2637) and TKO-bus Discriminator and Summing Module (DSM; HOSHIN 2548) which generates the trigger pulse and TDC inputs. The ADCs receive the divided signals within the time window of 100 nsec. Figure 4.8 shows a schematic diagram of the TKO front-end module and the discriminator and summing module (DSM). In the DSM, the signals are amplified and then divided into three signals. One is for the summing of 16 PMTs (Asum). The second one is fed to the updated discriminator. Figure 4.9 shows a schematic view of updated discriminator. The discriminated pulses go to the CAMAC multi-hit TDC modules which measure the arrival time and the pulse width. The timing resolution is 0.5nsec. This good time resolution is useful for the reduction of night sky background. The third one is fed into a non-updated discriminator. The discriminated signals ( 2 photoelectrons (p.e.) threshold) are fed into scaler circuit on this board. The scaler circuit counts the number of signals over the threshold during 700µsec when one-shot circuit is started by the external trigger every 15 sec. This value is called as scaler and is used for checking the night sky background and the electronic noises as is described in Chapter 6. This signal is also transformed into analog signal and is fed into summing-amplifier. This summed signal is called the logical sum (Lsum), and its pulse height is proportional to the number of hit PMTs. Using above two summed signal, the event trigger is made for the data acquisition (DAQ). Figure 4.10 shows a schematic diagram of the event trigger logic. Lsum signals from the DSMs are summed, and discriminated by the threshold corresponding to the number of hit PMTs (usually 3 hits). Asum signals of each AMP box are also discriminated ( 7 p.e.). The final event trigger is generated by the coincidence between the above two discriminated signals. The global DAQ trigger is made of the event trigger and the GPS trigger. The GPS trigger is made every 1 second for checking the time of DAQ system. The VME on-board CPU (FORCE 7V; Turbo Sparc 170MHz; Solaris 2.6) records the DSM scaler via VME-TKO interface, telescope tracking data via 100Mbase fast-ether network from the telescope control computer, and the data of weather 48

51 Figure 4.7: Block diagram of DAQ for the CANGAROO-II 10-m telescope. 49

52 Figure 4.8: Schematic diagram of the TKO front-end module and the discriminator and summing module. Figure 4.9: Updated discriminator. This discriminator makes a logic pulses when a pulse exceeds the threshold level. 50

53 Figure 4.10: Schematic diagram of the event trigger logic. and cloud monitor via RC-232C interface. All the data from the VME-bus (ADC, GPS) and CAMAC data (TDC, visual scaler, interrupt register) are recorded by Linux PC workstation (Intel Pentium-II 266MHz; Redhat 6.2; Linux kernel 2.2) via VME-PCI and CAMAC-ISA interfaces, respectively. The software used in DAQ system is the portable DAQ system UNIDAQ [66]. The typical data size of one event is 1.5 kbytes, and the data rate is 45kbyte/sec in average. The DAQ system can accept up to 80Hz triggers with a dead time of 20%. In order to check the quality of the data in real-time, another Linux PC workstation simultaneously receives the data and displays them. This computer also serves the system clock via NTP for all network computers by receiving GPS clock individually. 51

54 Chapter 5 Observations and Calibrations 5.1 Observations RX J was observed with the CANGAROO-II 10-m telescope in 2002 and The pointing direction was NW rim, (α, δ)= (8 h 48 m 59 s, ) (J2000), where the maximum X-ray emission was observed. RX J culminates at a zenith angle of 15. As was described in Chapter 3, the cosmic-ray can be eliminated using imaging air Cherenkov technique. After the selection, however, the cosmic-ray events still remain because of the large number of events compared with the gamma-ray events. Therefore we should take the background observation (OFF-source runs). The observation mode was Long ON/OFF. In this mode, the observations of the target (ON-source runs) are carried out during before and after the culmination for 1-5 hours typically. On the other hand, we take OFF-source runs with the offset along the axis of Right Ascension in order to take the same tracking as ON-source runs. This mode can maximize the ON-source exposure time. A total of 187 hours data was obtained. The observation times are summarized in Table 5.1. Observation Date T on (min) T off (min) 16-Dec Feb Jan Feb Total Table 5.1: Summary of the observation periods. 5.2 Calibrations The data were calibrated using a LED (Light Emitting Diode) light source located at the center of the 10-m reflector, 8-m from the camera [51]. A quantum-well type blue LED (NSPB510S, λ 470 nm, Nichia Corporation, Japan) was used. The camera was illuminated with an input pulse of 20 nsec width during the calibration runs before and after the observations. A light diffuser was placed in front of the LED in order to obtain a uniform yield on the focal plane. The data obtained by the calibration runs were calibrated with CALIB10 module in FULL [70]. 52

55 5.2.1 Terminologies Before starting to present the calibration and analysis methods, we define the terminologies. Table 5.2 summarizes the terminologies used for the analysis in this thesis. ADC TDC scaler ADC sum Pulse height of the signal for each PMT, i.e. proportional to the photon density. Arrival time of the signal pulse. Number of signals over the threshold during 700µsec. Sum of ADC for each PMT under some condition. This value is proportional to energy. Table 5.2: Terminologies used for the analysis in this thesis Field Flattening Each ADC value is subtracted by its pedestal value. The pedestal value was measured using the data of the calibration run, which was carried out without any external light. Two data sets of different luminosity of the LED, namely 0dB and 1dB LED-run data, were used for this calibration. At first, the mean value for each PMT (a suffix i implies the i-th PMT) Q i (xdb) and the averaged value for all PMTs Q ave (xdb) were obtained for x = 0dB and 1dB, respectively. In order to reduce the Night Sky Background (NSB), arrival timing of hit PMT was used. The signals of which timing were within 75 nsec were used. The relative gain value for the i-th pixel G i were obtained by fitting the linear function given as G i Q ave(xdb) Q i (xdb). (5.1) The obtained relative gain for each ADC is normalized to 1 by the mean of the relative gains of all PMTs Time-walk Corrections If the pulses are larger, the triggering timings are earlier. The correction of this effect is called as Time-walk corrections. Figure 5.1 shows the arrival timing of hit PMT (TDC) distributions before or after the Time-walk correction. The sharper peak was obtained. This helps us to discriminate shower data from NSB photons because the NSB photons come constantly and tend to make the flat TDC distribution Rejection of Bad Channels Bad channel means the PMTs which did not have any signal or hit rates of which were extremely high or low. The judgment conditions of bad channel are as follows: No signal is counted. 53

56 N (arbitrary unit) TDC count (nsec) Figure 5.1: Arrival time (TDC) distributions. Before the Time-walk corrections (dotted line). After the correction (solid line). TDC hit rate or ADC value shows the deviation five times larger in R.M.S., from its average value. Mean value of ADC exceeds the value of five times larger than the average for all pixels. Empty pixel at the corner of camera plane DST10 After the above calibration, the data were processed with DST10 module in FULL [22]. This module aborts the ADC data without TDC data. As shown in Fig. 5.2, such data are mainly low signals due to NSB photons ADC Conversion Factor In order to compare the observation and the simulation as is described in Subsection 6.5.1, a conversion factor from the ADC value to the photoelectron (p.e.) is required. Comparing the observation data with the Monte-Carlo simulations of protons about the rate and the relation between the total ADC counts and the total number of pixel hits, this factor was estimated to be [ADC ch/p.e.] by Itoh et al [48]. 54

57 N (arbitrary unit) ADC (count) Figure 5.2: ADC distributions for pixels, after pedestal subtraction, where N denotes the number of hits per 10-ADC count. The solid line is those with TDC hits and the dotted line without of them. 55

58 Chapter 6 Analysis The analysis of the data for RX J was carried out by the author. The analysis methods are presented in this chapter. 6.1 Reduction of the Night Sky Background (NSB) Generally photons emitted from air showers generated by gamma- and cosmic ray showers are detected with IACTs. These photons tend to form a cluster on the focal plane, and make a pulse with a width of a couple of 10 nsec. On the other hand, the Night Sky Background (NSB) photons tend to form separate images with low signal, and arrive randomly. To minimize the effects of the NSB, we selected events which formed clusters ADC Distributions Figure 5.2 shows a typical ADC distribution whose peak is located at ADC= 250. Therefore the threshold of ADC were selected as 300 in 2002 and 280 in 2003, respectively. The reason why the value in 2003 is lower is due to the difference in the hardware condition and the deteriorate of the mirror reflectivity, as is mentioned in Section 6.2. The thresholds of ADC are around 3.3 photoelectrons (p.e.) Clustering Using this ADC threshold, we looked into the good trigger cuts for the NSB reduction. Figure 6.1 shows the TDC distributions for various cluster sizes. Tna (Threshold n-adjacent) clusters mean the clusters with n adjacent pixels with signals above the threshold. We selected a T5a trigger cut TDC Cut Signals with TDC too far from the mean TDC were also rejected as NSB photons. Figure 6.2 shows the distribution of TDC after T5a-clustering and adjusting the mean TDC of each event to 0. We selected signals within ±50 nsec. 56

59 N x TDC count (nsec) Figure 6.1: TDC distributions for various cluster sizes. For all pixels (top). For pixels which satisfy T4a-clustering, and do not satisfy T5a (middle). For pixels which satisfy T5a-clustering and do not satisfy T6a (bottom). N denotes the number of hits per TDC count normalized to 1 nsec. N (arbitrary unit) TDC count (nsec) Figure 6.2: Distributions of TDC for a typical run after T5a-clustering and adjusting the mean TDC of each event to 0. Data inside the arrows were selected. The standard deviation is 13.2 nsec. 57

60 6.2 Cloud Cut and Elevation Cut As shown in Figure 6.3, the rates of the shower events decreased due to clouds (we can also use the information of the log books) and they also change due to the change of the energy threshold of cosmic rays according to the elevation. In order to satisfy constant N/ 5min Elevation(degree) Observation time(min) Figure 6.3: Change of event rate due to clouds and the change of elevation. Fine day (left). Cloudy day (right). Vertical axis is the number of events per 5 minutes (top) and elevation (bottom). acceptance and energy threshold, we do not use the data with low elevation or low event rate. Table 6.1 is a summary of the cloud cut and elevation cut conditions. The difference 16-Dec Jan Feb Feb Cloud cut (events/5min) Elevation cut (degree) Mean rate after cut (events/5min) Table 6.1: Cloud cut and elevation cut conditions and the resulting mean event rate. in the event rate between 2002 and 2003 was due to the difference in the hardware trigger condition and the deterioration of the mirror reflectivity. The trigger condition in 2003 was adjusted for stereo observation. Observation time after these selections is shown in Table % ON-source data and 73% OFF-source data remained. Figure 6.4 is shower rates as a function of the cosine of the zenith angle. Nearly constant event rates are achieved in each year. 58

61 Observation Date T ON (min) T OFF (min) T ON / T OFF 16-Dec Feb Jan Feb Total Table 6.2: Observation time after the pre-selection. Shower rate(hz) cos(zenith angle) Figure 6.4: Shower rates versus cosine of the zenith angle in 2002 and

62 6.3 Selection of Bad Pixels due to Starlights and Electrical Noises To monitor star light and electrical noise, the telescopes are equipped with scaler circuits. The hit rates within 700µsec of all pixels are monitored every 15 sec. The distributions of scaler counts for pixels obtained from and data are shown in Figure 6.5. Cut N (arbitrary unit) Scaler (count/700µsec) Figure 6.5: Distributions of scaler counts for pixels. N denotes the number of hits over threshold for 700 µsec. conditions of scaler counts are summarized in Table 6.3. The gain of PMTs were adjusted Observation date Mean (count) Cut (count) Cut ratio (%) 16-Dec Feb Jan Feb Table 6.3: Mean scaler counts, cut conditions, and cut ratios in 2002 and before the observation in 2003, therefore a slightly looser cut value was adopted. The bright star with the magnitude of 4.1, however, exists at 0.65 from the center. Figure 6.6 shows the optical images obtained by Digital Sky Survey. Figure 6.7 shows the maps with the correction of the rotation of the field. At scaler 15, star images can be seen in the ON-source data and noise in the OFF-source data. To investigate the effect of the stars further, we calculated the positions of the star s around the target. Figure 6.8 shows the tracks of the stars on the focal plane during the observation. Figure 6.9 shows the integral observation time distribution on the focal plane for the bright stars. The weights were calculated assuming the starlights have Lorentzian distributions and apparent luminosities of them follow Pogson s law. This figure indicates that the star with the magnitude of 4.1 makes a dominant effect. We, therefore, cut the pixels around 60

63 Figure 6.6: Optical image around the NW rim of RX J taken by Digital Sky Survey. Field of view is Figure 6.7: Scalar maps with the correction of the rotation of the field of view using the data in ON-source data with scaler 15 (upper left), scaler < 15 (lower left), OFF-source data with scaler 15 (upper right), scaler <