Preliminary measurement of cosmicray Carbon and Oxygen spectra with CREAMII

Transcription

1 Università degli Studi di Siena Facoltà di Scienze Matematiche Fisiche e Naturali Tesi di Dottorato in Fisica Sperimentale PhD Thesis in Experimental Physics XIX Ciclo Preliminary measurement of cosmicray Carbon and Oxygen spectra with CREAMII Riccardo Zei Relatore (Supervisor): Prof. Pier Simone Marrocchesi Siena, Maggio 2007

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3 Contents Introduction 1 1 Cosmic rays Energy spectrum and mass composition Acceleration Propagation General transport equations Solution for stable nuclei The LeakyBox Model The secondarytoprimary ratio Source composition Structures in the Energy Spectrum Measurements Techniques Indirect measurements Direct measurements The CREAM Detector Science objectives Instrument design Timing Charge & S3 Detectors Transition Radiation & Cherenkov Detectors Silicon Charge Detector Fiber Hodoscopes Target and Calorimeter Trigger System Balloncraft CREAMII ight and performance The second ight of CREAM Trigger logic Livetime Calorimeter performance High voltage and thermal behaviour Frontend electronics behaviour The Silicon Charge Detector performance Monte Carlo FLUKA: a brief introduction CREAMII Monte Carlo simulation i

4 ii Contents 4.3 Validation of Monte Carlo simulation Heavy ion simulation Shower reconstruction The event generation The energy interval The ux normalization The generation volume The Monte Carlo production Data Analysis Energy measurement with the Calorimeter Energy calibration LOW/MID ratios High voltage scale factor Observed energy deposit Charge Identication in the SCD Energy loss of charged nuclei in matter Principle of operation for the charge measurement Uniformity of SCD response Calibration of the silicon pixel response CREAM Analysis Software Data Analysis Tracking particles with the calorimeter Charge reconstruction in SCD Data sample Carbon and Oxygen Fluxes Formulation of the Unfolding problem Determination of the unfolding matrix Solution strategy Iterative unfolding Correction to the Top of atmosphere The Earth's atmosphere Monte Carlo simulation Absolute Flux determination The centering of data points Results and discussion Stability checks Conclusions 129 References 131

5 Introduction The present thesis summarizes my research activity, during a period of three years as a PhD student, in the Cosmic Ray Energetics And Mass (CREAM) experiment. This multimission balloon experiment is designed to provide direct measurements of the individual energy spectra and elemental composition of cosmicray nuclei at energies approaching 15 ev. Cosmic rays are charged (and neutral) particles coming from the space that span a very wide energy range (from MeV up to 20 ev). The main physics objectives driving the experimental study of the primary cosmic rays include the identication of their sources; a detailed understanding of the mechanism of their acceleration; the clarication of their propagation in the Galaxy and their interactions with the interstellar medium. To date, they are believed to be accelerated by remnants of supernova explosions in our Galaxy at energies below 15 ev per nucleon. At higher energies, their acceleration sites may be located in the Galaxy, whereas, at the end of the spectrum, an extragalactic origin seems to be necessary to explain the observed isotropy of the cosmic rays direction. A detailed overall comprehension is still missing. The CREAM experiment is a balloonborne mission aiming to test, in a relatively short time scale, the validity of a number of present astrophysical models. In particular, CREAM science objectives include the experimental test of models based on cosmicray acceleration by supernova shock waves, whose rigiditydependent limit could be reected in a composition change at energies 15 ev; the measurement of the energy dependence of the ratio of secondarytoprimary elements, that is closely related to the determination of the parameters of cosmic rays propagation in the Galaxy. CREAM completed two successful ights during a period of 13 months (from December 2004 to January 2006) from McMurdo, Antarctica, for a total of 70 days. A research group of INFN (Gruppo Collegato di Siena) and the Physics Department of the University of Siena partecipated in both the ight campaigns. The Italian group built a scintillator hodoscope (S2), for the rst ight, and a thin sampling tungsten/scintillating ber calorimeter, for the second ight. My contribution as a PhD student to the activity of the Siena group started in 2004, 1

6 2 Introduction when the construction of the calorimeter for the second ight was in progress. During the following three years, I worked on the analysis of the data collected during beam tests and the two ights. My work was partly dedicated to the development of a full Monte Carlo simulation of the CREAM instrument (FLUKA code based). At the same time, I studied and implemented new algorithms for the analysis and the reconstruction of events in the calorimeter. This information was then combined with the particle identication techniques of the other CREAM subdetectors to reconstruct the primary cosmicray spectra. This work presents the rst preliminary experimental results obtained analyzing the CREAMII ight data. They are also compared with the preliminary measurements of CREAMI, with two spacebased missions, HEAO and CRN, and two balloonborne experiments, TRACER and ATIC. These results have been approved by the CREAM collaboration to be presented at the International Cosmic Rays Conference 2007 in Merida, Messico. The outline of the thesis, organized in 6 chapters, is detailed as follows. In the rst chapter, a brief overview is given of the current knowledge about composition, acceleration and propagation of cosmic rays. At this point, the discussion is extended to the observational methods in astroparticle physics, focusing the attention on the direct detection with instruments above the atmosphere, on stratospheric balloons or in outer space. Chapter 2 illustrates the working principles of the CREAM instrument and describes its conguration during both ights. The preliminary analysis presented in this work is made using only the information provided by two subsystems, the SCD and the calorimeter. A summary of their ight performances is reported in the third chapter. To convert the energy deposit spectra into the primary energy spectra, it is necessary to unfold the instrument response. Detailed Monte Carlo simulations, needed for this process, play a prominent role in the analysis. Chapter 4 describes the CREAM instrument simulation, and the steps followed to assess its reliability. In the fth chapter, the analysis programs, developed by the author, as well as the Monte Carlo simulation, for the reconstruction and identication of the events in the CREAMII instrument, are described. At the end of the chapter, a selection of carbon and oxygen nuclei, reconstructed with the SCD and the calorimeter, is presented. In the sixth chapter, after describing the unfolding method used in the analysis, data are corrected for selection eciencies and for the losses due to interactions in the material up to the top of the instrument. The correction for the attenuation of the ux in the residual atmosphere is obtained by means of a dedicated Monte Carlo simulation and applied to the analyzed data. Finally, a preliminary measurement of cosmicray carbon and oxygen spectra, up to about 0 TeV, is presented. The stability of the tted spectral index is checked and the results are compared with previous measurements.

7 Chapter 1 Cosmic rays The year 1912 is usually considered as the year of the discovery of cosmic rays. Victor Hess [1] and later Werner Kolhörster [2] performed dedicated balloon ights in order to measure the altitude dependence of the natural radioactivity. The Kolhörster's ights reached the impressive height of 9 km. At the time this radiation was thought to originate from the Earth. However, in Hess' experiment the electroscope discharged more rapidly as the balloon ascended in the atmosphere which meant that the ionising radiation increased with altitude and therefore must originate from outside the Earth's atmosphere rather than from the ground. For some time it was believed that the radiation was electromagnetic in nature. It was only in 1926 that Millikan, at rst skeptical of Hess' work, used for the rst time the expression cosmic rays. Millikan showed also that the atmosphere does not produce cosmic rays, but acts only as a calorimeter. In 1929, an experiment made by Walter Bothe and Kolhörster based on coincidence measurements using two Geiger detectors allowed the charged nature of cosmic rays to be determined [3]. A few years later (1932) Carl Anderson used a cloud chamber and discovered the positron by measuring the curvature and energyloss of cosmicray particles in a magnetic eld [4]. This was the rst experimental evidence of the antimatter predicted by Dirac. That discovery was an important milestone in the development of particle physics. In the following years cosmic rays were employed as the only available lowcost source of high energy particles; they were soon replaced by the accelerators. In the last years, starting from the 1970s, cosmicray studies have been developed due to the growing interest in their origin and propagation, both related to the general knowledge of the universe evolution. 1.1 Energy spectrum and mass composition The energy spectrum of cosmic rays (fully ionized atomic nuclei) spans a wide range in energy. In Figure 1.1 the interval of energies from GeV to 20 ev is shown. Over these decades the ux decreases by about 30 orders of magnitude rather featureless, following roughly a power law dn/de E γ. The powerlaw behaviour indicates a nonthermal origin of the particles. To reveal small structures in the shape of the energy spectrum, the ux is usually multiplied by the energy to some power. The energy spectrum multiplied by E 3 is depicted in Figure 1.2. In this representation the spectrum looks rather at and ne structures can be recognized, indicating small changes in the spectral index γ. The most important are: 3

8 4 Cosmic rays -1 sr s GeV) 2 Flux φ (m (1 particle per m -second) Fluxes of Cosmic Rays -12 Knee 2 (1 particle per m -year) Ankle 2 (1 particle per km -year) Energy (ev) Figure 1.1: Allparticle cosmicray spectrum

9 Cosmic rays 5 Flux dφ/de 0 E 0 3. [m -2 sr -1 s -1 GeV 1.5 ] AGASA + Akeno 20 km 2 + Akeno 1 km 2 AUGER BLANCA CASA-MIA DICE BASJE-MAS EAS-Top Fly's Eye Haverah Park Haverah Park Fe Haverah Park p HEGRA HiRes-I HiRes-II HiRes/MIA KASCADE (e/m QGSJET) KASCADE (e/m SIBYLL) KASCADE (h/m) KASCADE (nn) MSU Mt. Norikura SUGAR Tibet ASγ Tibet ASγ-III Tunka extragal. direct: galactic JACEE RUNJOB SOKOL Grigorov knee 2nd knee ankle Energy E 0 [GeV] Figure 1.2: Allparticle energy spectrum of cosmic rays, the ux is multiplied by E 3. The lines represent spectra for elemental groups (with nuclear charge numbers Z as indicated) according to the polygonato model [18]. The sum of all elements (galactic) and a presumably extragalactic component are shown as well. The dashed line indicates the average all-particle ux at high energies. the knee at E k 4.5 PeV where the powerlaw spectral index changes from γ 2.7 at low energies to γ 3.1 the second knee at E 2nd 400 PeV 92 E k, where the spectrum exhibits a second steepening to γ 3.3 the ankle at about 19 ev: above this energy the spectrum seems to atten again to γ 2.7. The origin of these structures is expected to be the key element in the understanding of the origin of cosmic rays (CRs). At energies around 1 GeV/n all elements known from the periodic table with nuclear charge number Z from 1 to 92 have been found in CRs [5]. The relative abundance of elements in CRs is very similar to the abundance found in the solar system, which indicates that CRs are regular matter, but accelerated to very high energies. This is illustrated in Figure 1.3(a)1.3(b). Both compositions show a peak at the elements C, N, O and Fe, and a clear evenodd eect in the relative stabilities of the nuclei (i.e. nuclei with an even atomic number are more abundant than those with an odd one). However, even if a general good agreement between cosmic rays and solar matter can be stated, there are distinguished dierences as well. In fact, the two groups of elements Li, Be, B and Sc, Ti, V are almost absent in the solar system while a signicant amount of them is found in cosmic rays. Same consideration can be applied to the elements below the lead group (Z = 82). The abundances observed in cosmic radiation are interpreted as a result of the collisions of nuclei with the interstellar medium (ISM) during the propagation of cosmic rays from their sources to the Earth and the subsequent fragmentation of the primary nucleus into

10 6 Cosmic rays Relative abundance Li Be B C N O Ti V Sc Fe Charge number Relative abundance (Fe 1) ARIEL 6 Fowler HEAO 3 UHCRE SKYLAB TIGER Trek MIR Tueller + Israel sol. syst Nuclear charge Z (a) (b) Figure 1.3: (a) Comparison between the cosmic rays and the Solar System element composition, both given relative to carbon. Filled circles: cosmic rays. Open circles: Solar System elements. (b) Relative abundance of cosmic rays (Z > 28) normalized to Fe 1 from various experiments. lighter elements. The products of such nuclear fragmentation processes, like boron nuclei, are called secondary cosmic rays, while elements, like carbon, that mostly originate from nucleosynthesis in stars are referred to as primary cosmic rays. Therefore the relative abundances of secondary nuclei with respect to their progenitors give a measure of the (average) amount of material that cosmic rays typically traverse between injection and observation. Collisions of primary cosmicray nuclei with the ISM also cause production of long lived radioactive isotopes, which can be used as cosmicray clocks. By comparing the observed abundances of these radioactive nuclei with the amount of stable secondary species in the cosmic radiation, one can establish for how long cosmic rays are conned in the Galaxy and investigate the density distribution of the gas encountered by the particles. In fact, since secondary radioactive cosmicray species are created and decay only during propagation, their steadystate abundances are sensitive to the connement time, if their mean lifetimes are comparable or shorter than this time. It has been deduced that an average CR spends several million years wandering around our Galaxy before reaching the Earth. Since CRs are charged particles, most of them, while propagating in the Galaxy, have an energy low enough to be deected by the interstellar magnetic eld; therefore, once they reach the Earth, all directional information about their source has been lost. 1.2 Acceleration The bulk of CRs is assumed to be accelerated in blast waves of supernova remnants (SNRs). This goes back to an idea by Baade and Zwicky who proposed SNRs as cosmic ray sources due to energy balance considerations [7]. They realized that the power necessary to sustain the cosmicray ux could be provided when a small fraction % of the kinetical energy released in supernova explosions is converted into CRs. In 1949 Fermi

11 Cosmic rays 7 proposed a mechanism to accelerate particles with moving magnetic clouds [8]. This led to todays picture that the particles are accelerated at strong shockfronts in SNRs through rst order Fermi acceleration. This theory predicts spectra at the sources following a power law dn de E 2.1 (1.1) Power The energy density of cosmic rays, that is locally observed, amounts to about ρ cr 1 ev/cm 3. The power required to sustain a constant cosmicray intensity can be estimated as P cr = ρ crv 5 33 J/s τ esc where τ esc 6 6 years is the connement time of cosmic rays in a containement volume V 4 66 cm 3 (the Galaxy or the galactic halo). The observed rate of supernova explosions in our Galaxy is approximately 23 per century, which yields a power of around 3 35 J/s. Therefore, only a few percent of the total energy released in a typical explosion needs to be transferred to CRs to maintain the observed intensity. Fermi mechanism Let us suppose that a charged particle of energy E 0 enters a magnetized region R A of the interstellar medium. This particle crosses n times this region before it escapes denitely, each time gaining an energy fraction ξ = E/E, with ξ > 0. After n acceleration cycles, the energy acquired by the particle is E n = E 0 (1 + ξ) n (1.2) From eq. (1.2), the number of encounters needed to reach a certain energy E is given by ( E ) n = ln / ln(1 + ξ) (1.3) E 0 Let us indicate with P esc the probability of the particle to escape from R A after each acceleration cycle. The quantity (1 P esc ) n is indeed the probability of the particle to remain in the acceleration region after n cycles. The number of particles with energies greater than E, after n encounters, is N( E) (1 P esc ) m = (1 P esc) n (1.4) P esc m=n Substituting n in this expression yields N( E) 1 P esc ( E E 0 ) α (1.5)

12 8 Cosmic rays upstream downstream Ε=Ε+ Ε E β= v/c 1 E 2 - u 1 V = -u 1 + u 2 (a) (b) Figure 1.4: (a) Fermi original idea: charged particle scatters inside the cloud of magnetized plasma.(b) Fermi acceleration at a shock wave front. The strong wave propagates with velocity u 1 and the shocked gas ows away from the shockfront with velocity V = u 1 + u 2. The thick vertical line represents the shock wave front, the thin line describes the path of a particle. where α = ( ln ) 1 1 P esc ln(1 + ξ) P esc ξ = 1 ξ T cycle T esc (1.6) Here α only depends on the fractional energy gain ξ and on the escape probability P esc, which can be expressed as the ratio of the characteristic time for the acceleration cycle T cycle over the characteristic escape time T esc. As can be seen from eq. (1.6), Fermi found that this acceleration mechanism leads to a powerlaw spectrum of particle energies. In fact, dierentiating eq. (1.5) with respect to energy E, the resulting dierential energy spectrum is dn(e) de E (α+1) (1.7) For shockwave acceleration [11], it turns out that α 1.1, so that the dierential spectrum index is 2.1, compared with the observed value of γ = 2.7. The steeper observed spectrum can be accounted for if the escape probability P esc is energy dependent. A schematic picture of the Fermi idea is shown in Figure 1.4(a). When cosmicray particles collide with moving clouds of plasma, they can either gain or lose energy in the scattering, however, it could be shown [9] that on average there is a net energy gain E : E 4 3 E 0β 2 where β is the plasma velocity with respect to the speed of light. As the average energy gained is proportional to β 2, this mechanism is called second order Fermi acceleration.

13 Cosmic rays 9 Though this mechanism predicts a power law for the dierential energy spectrum of cosmic rays, because of the quadratic dependence, the process is quite slow and inecient in comparison to what is needed to sustain the observed cosmicray spectrum. It was shown [11] that the Fermi idea can be modied in order to describe more powerful acceleration which takes place in supernovae shocks. In the case of an innite plane shock front, which moves with velocity u 1, the shocked gas has a velocity v = u 1 + u 2 relative to the unshocked gas (upstream), with u 2 < u 1. In this scenario, an acceleration cycle is characterized in the following way: an encounter begins when a charged cosmicray particle crosses the shock front from the upstream region under an angle θ 2 against the direction of the front the particle scatters on the irregularities in the turbulent magnetic eld, which is carried along with the moving plasma in the downstream region the encounter ends when the particle crosses the front again, this time from the downstream to the upstream region. The situation is illustrated in Figure 1.4(b). With this denition, every encounter results in an energy gain, which can be calculated as described below. If the initial energy of the particle is E 1 pc, i.e. the particle is suciently relativistic, a Lorentz transformation gives its energy in the rest frame of the moving plasma in the following form E 2 = γe 1 (1 β cos θ 2 ) (1.8) Energy is transferred to the particle when it collides with the fast moving magnetic eld that is carried with the plasma. In the rest frame of the moving plasma, the scatterings are not due to actual collisions, but are caused only by the motion in the magnetic eld and are therefore elastic. Thus the energy of the particle in the moving frame remains unchanged until it leaves the downstream region. Then its energy in the laboratory frame is E 3 = γe 2 (1 + β cos θ 1 ) (1.9) The fractional energy gain from a single encounter can then be calculated as E E 1 = E 3 E 1 E 1 = 1 β cos θ 2 + β cos θ 1 β 2 cos θ 2 cos θ 1 1 β 2 1 (1.) Assuming an isotropic ux, it turns out [] that the probability of the particle crossing the shock is proportional to sin θ cos θdθ, i.e. cos θ 1 = π/2 0 2cos 2 θ sin θdθ = 2 3 (1.11) and cos θ 2 = 2/3, which leads to an average fractional energy gain E E 1 = β β2 1 β 2 1 (1.12)

14 Cosmic rays As long as the shock velocities are nonrelativistic, this can be approximated as E E β (1.13) The average fractional energy gain thus depends only on the velocity of the shocked gas relative to the unshocked region. Due to the linear velocity dependence, acceleration on shock wave fronts is referred to as rst order Fermi acceleration. Thus, both (second and rst order) mechanisms consist in a pair of crossings: in and out of the cloud, and back and forth across the shock. The rst order Fermi acceleration is believed to be at work in many dierent types of shocks: the termination shock of the Solar and the galactic winds, the acceleration shock near a supermassive black hole (which is believed to exist in the center of our Galaxy) and the expanding envelopes of exploded supernovae. Supernova blast waves As shown above, diusive rstorder shock acceleration works by virtue of the fact that particles gain an amount of energy E E at each cycle, when a cycle consists of a particle passing from the upstream (unshocked) region to the downstream region and back. At each cycle, there is a probability that the particle is lost downstream and does not return to the shock. Higher energy particles are those that remain longer in the vicinity of the shock. In fact, the limited time during which the supernova blast wave is active determines the maximum energy per particle that can be achieved by this mechanism. After a time T the maximum energy attained is E max Zeβ s BT V s (1.14) where β s = V s /c refers to the velocity of the shock. This results in an upper limit, assuming a minimal diusion length equal to the Larmor radius of a particle of charge Ze in the magnetic elds B behind and ahead of the shock. Using typical values of Type II supernovae exploding in an average interstellar medium (T 00 years and B ISM 3 µg) yields [51] E max Z 0 TeV (1.15) More recent estimates give a maximum energy up to one order of magnitude larger for some types of supernovae [22] E max Z 5 PeV (1.16) As the maximum energy depends on the charge Z, heavier nuclei can be accelerated to higher energies. In this model, this would imply the existence of consecutive cutos in the energy spectra for individual elements proportional to their charge Z, starting with the proton component. As previously observed, the overall similarity between cosmic rays and matter in the solar system (see Figure 1.3(a)) indicates that CRs are regular matter, but accelerated to extreme energies. The supernova acceleration model works well below the knee, but it cannot produce enough CRs above 18 ev. Moreover, since the gyroradius r L (or Larmor radius) in the

15 Cosmic rays E max~ βzbl 9 Protons β=1/300 Protons β=1 log(magnetic field, gauss) 3 3 Neutron star White dwarf Fe Crab AGN SNR x RG lobes Colliding galaxies Galactic disk halo x Virgo 9 Clusters au 1 pc 1 kpc 1 Mpc log(size, km) Figure 1.5: The Hillas diagram. Magnetic eld strengths and characteristic dimensions are plotted for various candidate sources of accelerated particles. Objects below the blue dashed diagonal line cannot accelerate protons to 20 ev. galactic magnetic eld at 18 ev is of the order of the size of the galactic disk, a transition from galactic to extragalactic cosmic rays is expected to occur around this energy. Hillas Condition In addition to the acceleration mechanism, a simple argument can be used to express the upperbound energy which can be reached in the acceleration process. The Larmor radius r L of a relativistic particle of charge Ze in a magnetic eld B (strictly the component of

16 12 Cosmic rays B normal to the particle velocity) is r L = p ZB E ZB (1.17) where E is the particle energy in units of 15 ev, B is in µg and r L in parsecs. Supposing that there is a region of acceleration with characteristic dimension R: the idea is that a charged particle can be accelerated inside such a region up to an energy where its Larmor radius r L reaches the critical value R/2, i.e. 2r L R. This condition leads to the Hillas inequality [14]: B R > 2 E (1.18) Z This formula (1.18) puts a constraint on the combination of source size R and magnetic eld B for a certain maximum energy. Moreover, eq. (1.18) can be stated in terms of the maximum energy that a charged particle can be accelerated to, as: E = βzb R 2 (1.19) where the factor β ( 1) is the fraction of the maximum energy that can be attained within the region R and is related to the details of the acceleration mechanisms at the particular site (in simple shock acceleration β = β s = V s /c where V s is the velocity of the shock). This scenario can be expressed in terms of the socalled Hillas plot which is illustrated in Figure 1.5. Various candidates of cosmic accelerators are placed in the plot: neutron stars, white dwarfs (both are representants of direct acceleration), supernova remnants, active galactic nuclei (AGN), lobes of radio galaxies and colliding galaxies. It is evident that there are only a few possible sites for acceleration of cosmic rays up to the very end of the observable spectrum. Candidates lying below the diagonal line fail to satisfy condition (1.18): in Figure 1.5 the cases for 20 ev protons (blue dashed diagonal line) and 20 ev iron nuclei (green continuous diagonal line) with β = 1 (describing Fermi acceleration at an ultrarelativistic shock front) are plotted. For smaller β the resulting lines would lie higher. Clearly, the plot shows that, in order to accelerate charged particles to extremely high energies, either highly condensed objects with huge B or enormously extended sites are necessary. In either case, very high speeds are required. Among the excluded sites are supernova remnant envelopes. 1.3 Propagation After acceleration, CRs propagate through the Galaxy, being deected many times by the randomly oriented magnetic elds B ISM and experiencing several physical processes: diusion along magnetic eld lines, scattering on magnetic eld irregularities, nuclear fragmentation reactions with the interstellar gas and radioactive decays. Cosmicray propagation models take into account all these processes and attempt to explain the uxes observed on the Earth.

17 Cosmic rays General transport equations A general transport equation has been proposed by Ginzburg and Syrovatskij [ 16]: N i t = (D N i ) E [b in i ] un i + Q i p i N i + N k p k i (1.20) k>i where N i = N i (E, r, t) represents the density of cosmicray nuclei of type i at position r with energy between E and E + de. The rst term on the right hand side describes the diusion of particles; the diusion coecient D can be written as D = 1 3 λ Dv (1.21) where v is the particle velocity and λ D the diusion mean free path. The second term represents the gain of energy or loss (due, for example, to ionization), that is b i de dt (1.22) the mean rate at which the single nucleus i changes its energy. The convection term is the third one and represents convection with velocity u. For the following discussion it will be assumed that particles originate in pointlike sources whose intensity, spectrum and space time distribution are described by the functions Q i = Q i (E, r, t), where i denotes the nucleus type. The quantity Q i expresses the number of particles injected in the interstellar medium, per time unit and volume, in the energy range between E and E + de. The fth term represents losses due to collisions and decay, whose probability per time unit is p i = nβcσ i + 1 (1.23) γτ i where τ i is the lifetime of nucleus i for radioactive decay, β is the velocity of the particle, γ is the Lorentz factor, the interstellar gas is assumed to be hydrogen with density n n(r) (typical density n 6 m 3 ) and σ i is the inelastic crosssection. The last term describes the production of nuclei of type i from the interactions of nuclei of a dierent type (in the sum only heavier nuclei are considered). Relation (1.23) can be used to express p k i. The partial crosssection for the production of the nucleus i from the breakup of a nucleus of type k is σ k i and τ k i is the lifetime of the nucleus k with respect to the nuclear decay channel that leads to the production of a nucleus of type i. The production term can be written in the form of relation (1.23) because in fragmentation reactions of relativistic nuclei the kinetic energy per nucleon is nearly (to rst approximation) conserved. This process is known as spallation. The dierent models of the Galaxy can be represented using equation (1.20) with specic boundary conditions and approximations. The general approach is, rst, to get the parameters of the theory from the observed secondary abundances. In fact, in this case the source term is Q i = 0, so that the secondary cosmicray density is only related to the propagation parameters. Once these parameters are known, one can infer the composition at the sources from the observed primary abundances.

18 14 Cosmic rays Solution for stable nuclei The aim of this paragraph is to discuss the role of the stable secondary cosmic rays as tools to study the propagation mechanism within the Galaxy. The discussion is therefore limited to stable nuclei, setting the decay term in eq. (1.23) to zero. A stationary picture of the cosmicray propagation is taken into consideration, since several experimental data [] suggest that cosmicray intensity did not change significantly over the last hundred million years. As a consequence, the time dependence of eq. (1.20) vanishes. Moreover, the discussion is restricted to an energy range above several GeV/nucleon. In this way, transport equations are simplied since ionization losses can be ignored. With the above mentioned conditions eq. (1.20) simplies as follows: (D N i ) + nβcσ i N i = Q i + k>i nβcσ k i N k (1.24) Let us assume that the spatial distribution of sources is independent of the kind of nuclei, e.g. Q i (r) = q i χ(r) where q i are constants that determine the abundances of the various nuclei at the sources. It can be demonstrated [9] that eq. (1.24) has a solution of the form: N i (r) = N (σ) i (x)g(r, x)dx (1.25) 0 where N (σ) i and G, which are both function of the parameter x (expressed in g/cm 2 ), satisfy the following equations: dn (σ) i dx mnβc G x + σ i m N (σ) i k>i (D G) = 0 with mnβcg(r, 0) = χ(r) (1.26) σ k i m N (σ) k = 0 with N (σ) i (0) = q i (1.27) where m is the hydrogen mass. The function G should satisfy the boundary conditions for N i (r) related to the geometry of the containment volume. By espressing the transport equation in this form, it is possible to separate the strictly astrophysical aspects of the cosmicray propagation, related to the function G, from the aspects concerning the fragmentation by nuclear interactions with the interstellar medium,. In fact, given a specic composition q i at the sources, the quantities N (σ) i (x) represent the variation of the cosmicray composition due only to fragmentation processes after traversing an amount x of matter, and they are related only to the nuclear crosssections (σ i,σ k i ). The quantity G(r, x) is instead related only to the parameters of the considered propagation model: the diusion coecient ( D), the geometry of the containment volume, the source and interstellar medium distributions (n(r),χ(r)). related to the function N (σ) i The function G is usually called pathlength distribution. The interpretation of G(r, x) as pathlength distribution is clear from eq. (1.25), which states that the cosmicray density at the position r is obtained by averaging the composition expressed by N (σ) i (x)

19 Cosmic rays 15 over the distribution of path lengths G(r, x) travelled by CRs from their sources to the point r. For any propagation model that satises the requirement Q i (r) = q i χ(r) and can be expressed by a set of transport equations of the kind (1.24), all the information concerning the propagation processes is included in the function G(r, x). The following paragraph will show that the cosmicray secondarytoprimary ratio provides in principle direct information on the pathlength distribution and, as a consequence, on the cosmicray propagation mechanism The LeakyBox Model One of the more successful Galaxy models is the LeakyBox Model (LBM) and its variants. Within the framework of the LBM it is assumed that the density of the CRs and the distribution of the ISM and of the sources in the whole system (the Galaxy) are constant and uniform. The diusion term in eq. (1.24) is thus replaced by assuming a nite escape probability 1/τ esc from the connement volume: (D N i ) N i τ esc (1.28) In the scenario generally associated with the LBM, the Galaxy is described as a cylinder with radius 8 kpc and thickness pc; CRs are assumed to propagate freely inside the containment volume, where a magnetic eld of 36 µg exists. When CRs reach the cylinder boundary they bounce elastically, but there is a nite probability (increasing with the particle momentum) that they cross the boundary and escape from the Galaxy. The LBM set of transport equations which describes the propagation of stable nuclei is obtained by substituting the escape term in eq. (1.24) and dividing by the term mnβc: Q i mnβc N i λ i N i λ esc + k>i N k λ k i = 0 (1.29) In eq. (1.29) we have introduced the interaction lengths λ = m/σ, expressed in g/cm 2, where σ is the crosssection for the considered process. With these approximations the cosmicray propagation is described, for the various kinds of nuclei i, by a set of algebric equations which can be solved analytically. In eq. (1.29) the only parameter of the model is expressed as λ esc = mnβcτ esc (1.30) which represents the mean amount of matter, in g/cm 2, traversed by CRs before their escape from the connement volume. Such interpretation of λ esc is evident from the LBM pathlength distribution function. In fact, the solution of eq. (1.26), after the replacement of the diusion term with the escape term, gives as a result: G LBM (x) e x/λesc (1.31) One of the main constraints on the LBM parameter λ esc comes from the ratio of stable secondarytoprimary abundances. For semplicity sake, let it be taken into consideration a secondary nucleus of type S whose only parent nucleus is of type P. In this case the

20 16 Cosmic rays set of transport equations is reduced to a system of two equations: Q P mnβc N P λ P N S λ S N S λ esc + N P = 0 λ esc (1.32) N P = 0 λ P S (1.33) The solution of the above system gives the following expression for the secondaryto primary ratio: S/P = N S = λ esc/λ P S N P (1 + λ esc /λ S ) λ esc (1.34) λ S λ esc λ P S In eq. (1.34) it is clear that, in the extreme case of a secondary nucleus with a long interaction mean free path, the secondarytoprimary ratio S/P is directly proportional to the escape mean free path of CRs from the Galaxy The secondarytoprimary ratio The main results concerning the stable nuclear component of CRs within the framework of the LBM are summarized in the following. The observed abundances at a few GeV/n of the stable secondary nuclei Li, Be, B and Sc, Ti, V (see Figure 1.3(a)), produced by the spallation of the nuclei C, N, O and Fe respectively, are well described by assuming an escape mean free path from the connement volume of λ esc = 5 g/cm 2 (1.35) of equivalent hydrogen. Therefore it can be inferred that cosmic rays travel distances through the galaxy a factor thousand larger than the thickness of the galaxy. The energy dependence of the secondarytoprimary ratio (see B/C and (Sc+V)/Fe in Figure 1.6(a)1.6(b)) can be explained assuming an escape path length for particles with rigidity R and velocity β (w.r.t. light speed) of the form: λ esc (R) = ( βr 1.0 GV 26.7βg/cm 2 ) δ ( + βr 1.4 GV ) 1.4 (1.36) with δ [47]. The observed ratio of boron to carbon suggests that higher energetic particles traverse less matter as they escape earlier from the galaxy. Furthermore, it implies that the acceleration occurs before the propagation, else the ratio would be constant or increase with energy. At high energies, the pathlength according to relation (1.36) decreases as λ esc (R) = λ 0 ( R R 0 ) δ (1.37) with typical values λ 0 15 g/cm 2, and rigidity R 0 4 GV[17]. In this scenario the spectra observed on the Earth should be steeper as compared to the source,

21 Cosmic rays 17 B to C ratio HEAO-3-C2 ACE/CRIS CRN (Sc+V) to Fe ratio 0.1 HEAO-3-C2 ACE/CRIS Energy (GeV/n) Energy (GeV/n) (a) (b) Figure 1.6: Abundance ratio of boron to carbon (a) and scandium+vanadium to iron (b) in cosmic rays as function of energy. i.e. the spectral index γ should be smaller by the value of δ. Since the allparticle spectrum has γ 2.7, this means that the source spectrum is E 2.1, compatible with the spectral shape that is expected whether the CRs are accelerated by a rst order Fermi mechanism. Energy spectra of individual elements have been measured up to energies of about 14 ev by experiments above the atmosphere, the results being well compatible with power laws [5]. The spallation processes during the cosmicray propagation also inuence the shape of the spectra. It is assumed that the energy spectra of all elements have the same spectral index at the source. Taking the pathlength of the particles in the Galaxy into account, it is found that on the Earth the spectra of heavy nuclei should be atter as compared to light elements [5]. This fact is supported by direct measurements of individual energy spectra [5], e.g. the values for protons γ p = 2.71 ± 0.02 and iron γ Fe = 2.59 ± 0.06 dier as expected. During propagation processes, the regular component of the galactic magnetic eld will cause particles with charge Z to describe helical trajectories with a Larmor radius r L while the random eld component causes diusive propagation. With increasing energy (or momentum) it becomes more dicult to magnetically conne the particles to the Galaxy. Since r L 1/Z, it is expected that leakage from the Galaxy occurs for light elements (low Z) earlier as compared to heavy nuclei, i.e. protons leak rst and subsequently all other elements start to escape from the Galaxy. 1.4 Source composition Besides the source spectra, the elemental composition of the source at high energies can be directly studied. This can be done for a wide range of elements since those which

22 18 Cosmic rays (a) GCRS/LG (Fe=1) 1-1 preliminary 0 GeV/n GCRS/LG (Fe=1) 1 preliminary 00 GeV/n GCRS/LG (Fe=1) GeV/n First ionization potential [ev] Condensation temperature [K] (b) Figure 1.7: (a) Elemental abundance ratios (relative to Si) versus FIP. The circles show the ratio of the galactic cosmicray source (GCRS) to local galactic abundances. The triangles show the ratios of elemental abundances in solar energetic particle events to solar photospheric abundances. (b) Abundances of elements in the cosmicray source relative to local galactic abundances are plotted versus the FIP or versus the condensation temperature [46]. Data from TRACER [38] (squares) and CRN [41] (open circles). For reference see [20] and [46] have a too poor statistical precision for accurate spectral measurements can still provide useful abundances at high energies. In Figure 1.7(a) the cosmicray source abundances are compared with local galactic abundances. Also shown are the ratios of the elemental abundances in the solar wind at 512 MeV/n compared to the solar photospheric abundances. These data are plotted versus the atomic rst ionization potential (FIP). Ratios of individual cosmicray elemental abundances to the corresponding solar system abundances seem to be ordered by the FIP of the elements. An average depletion of elements

23 Cosmic rays 19 with higher FIP is clearly shown: in the GeV energy range elements with a FIP below ev are about 5 times more abundant in the cosmicray sources than are elements with a higher FIP [44]. It is interesting to compare the FIP eect in CRs with that seen with solar energetic particles. There are some general similarities between these two ratios which may suggest that energetic particles owing from stars could form a suprathermal seed population of particles which is subsequently accelerated as CRs. However, the H and He abundances are clearly underabundant relative to the solar values. The fact that solar values tend to lie close to the trend of the FIP established for the heavier nuclei, suggests that the cosmicray source material comes from regions which are enriched in heavier nuclei, or that some other mechanism suppresses the initial acceleration of the lighest nuclei. It has also been found that the abundance ratios of galactic cosmic rays to the solar system values scale with the condensation temperature T C. Refractory elements (T C > 1250 K) are more abundant than volatile elements (T C < 875 K). A model which could perhaps account for the FIP eect would have a source ejecting its outer envelope for a long period with a FIP selection eect (as the sun does) followed by a supernova shock which sweeps up and accelerates this material. It has been proposed that the FIP eect is due to nonvolatiles being accelerated as graines, while the volatiles are accelerated as individual nuclei [45]. Recent measurements of the TRACER experiment allow to investigate these eect at higher energies, namely at 0 GeV/n and 1 TeV/n [46]. The decreasing ratio as function of FIP and the increase as function of condensation temperature, known from lower energies is also pronounced at energies as high as 1 TeV/n, as can be inferred from Figure 1.7(b). 1.5 Structures in the Energy Spectrum Many possible origins for the knee are discussed in the literature [19]. The most popular models assume a nite energy attained during the acceleration process and leakage from the Galaxy as discussed previously. In both scenarios, the energy spectra of elements exhibit a cuto at an energy proportional to the nuclear charge Z and the knee in the allparticle spectrum is caused by the cuto of light elements, starting with protons. All other elements follow subsequently and above a certain energy no more particles are left. On the other hand, the measured allparticle ux extends up to 20 ev. The highest energy particles are usually believed to be of extragalactic origin because the Larmor radius of a proton with an energy of 20 ev in the galactic magnetic eld is r L 36 kpc, comparable to the diameter of the Galaxy. The transition region from galactic to extragalactic CRs is of particular interest: key features are the origin of the second knee and ankle. Transition from galactic to extragalactic cosmic rays The location in energy of a transition from galactic to extragalactic cosmic rays is a long standing question. The main lines of thought in this respect are briey summarized below. It has long been thought that the ankle, at (0.5 1) 19 ev, is a feature arising

24 20 Cosmic rays from the intersection of a steeply falling galactic spectrum of cosmic rays with a atter spectrum of extragalactic cosmic rays. In the socalled ankle scenario the extragalactic component has a very at generation spectrum E 2 which naturally intersects the steep ( E 3.1 ) galactic component. This view implies that very ecient acceleration processes are needed in order to extend the galactic cosmicray spectrum to energies 19 ev. Moreover, the dip seen in the spectrum between 18 ev and 19 ev (see Figure 1.2) is interpreted as the signature of a possible interplay between galactic and extragalactic spectra. Reviewing the properties of CRs accelerated in SNRs, Hillas nds that a conservative estimate of the maximum energy achieved during the acceleration in supernova remnants is not sucient to explain the allparticle ux up to 17 ev. In this scenario a second (galactic) component, accelerated by some type II supernovae, is necessary to explain the observed ux at energies above 16 ev [15]. In this way, a contribution of extragalactic particles is relevant for energies higher than 17 ev. Another possibility is a signicant contribution of ultraheavy elements (heavier than iron) to the allparticle ux at energies around 400 PeV [18]. In this approach the second knee is caused by the fallo of the heaviest elements with Z up to 92. It is remarkable that the second knee occurs at E 2nd 92 E k, the latter being the energy of the rst knee. In this way, the transition to extragalactic cosmic rays takes place at energies higher than 5 17 ev. In the socalled dip scenario, the dip seen in the spectrum between 18 ev and 19 ev is proposed [21] to be caused by the combination of adiabatic losses (expansion of the universe) and electronpositron pair production of CRs on cosmic microwave background (CMB) photons p + γ 2.7K p + e + + e This dip was shown to t very well the observed spectra in AGASA, HiRes, Yakutsk and Auger experiments, as shown in Figure 1.8. As a consequence of this interpretation, a very important point to be made is that the position of the dip is independent of astrophysical details and is xed by the rates of adiabatic losses and pair production. In this scenario the transition between galactic and extragalactic cosmic rays is described as the intersection of a steep galactic spectrum at E > ev with a at extragalactic proton spectrum at E < 1 18 ev. The low energy part of the dip ts the second knee ( ev) which is interpreted as the signature of the transition to an extragalactic population. Below the second knee the predicted extragalactic spectrum attens and drops below the ux of galactic cosmic rays. Both models make clear predictions: as far as the chemical composition is concerned, the two models dier mostly. In fact, in the ankle scenario, the galactic CRs are expected to be mainly iron nuclei (see Figure 1.5) for energies higher than 18 ev, while the dip scenario predicts a strong dominance of protons. It is important to remark that the estimates of the various approaches are quite similar. The dierences are caused by dierent estimates of the tail of the galactic ux to highest

25 Cosmic rays modification factor -1 Akeno-AGASA η ee modification factor -1 HiRes I - HiRes II η ee -2 γ g =2.7 η total -2 γ g =2.7 η total (a) E, ev (b) E, ev 0 0 modification factor -1 Yakutsk η ee modification factor -1 Auger η ee -2 γ g =2.7 η total -2 γ g =2.7 η total (c) E, ev (d) E, ev Figure 1.8: Predicted dip in comparison with (a) AGASA [30], (b) HiRes [31], (c) Yakutsk [28] and (d) AUGER [33] data. For reference see [23] energies. New measurements of the mass composition in the energy region of the second knee will help to discriminate between the dierent scenarios. GreisenZatsepinKuzmin cuto At energies around 20 ev, a crucial issue is represented by the measurement of the socalled GZK suppression eect. Soon after the discovery of 2.7 K black body relict radiation, it was pointed out by Greisen [12], Zatsepin and Kuzmin [13] that very high energy protons react with the relic microwave photons p + γ 2.7 K + p + π 0 (1.38)

26 22 Cosmic rays Figure 1.9: Mean energy of protons as a function of propagation distance. For reference see [24] The reaction threshold is about 6 19 ev. The process (1.38) is often called photo disintegration. Below the threshold energy the proton attenuation length is 00 Mpc, while above the threshold it is reduced only to 20 Mpc. Similar eect appears also for heavier nuclei. If the ultra high energy sources are extragalactic, the primary particle has to pass several subsequent interactions of the type (1.38) until it reaches the Earth. In each of these interactions the energy of the incident particle is reduced, and the process stops when the resulting energy drops below the threshold of photodisintegration. The eect is depicted in Figure 1.9 where the mean proton energy is plotted as a function of the propagation distance. An endpoint of the cosmicray spectrum is thus expected around 20 ev. The present experimental situation is not clear as the AGASA and HiRes spectra are in mild contradiction with each other, which may well be due to statistical uctuations and a possible systematic error in the energy determination with the two dierent techniques used. The AUGER experiment should soon settle this issue. 1.6 Measurements Techniques In order to clarify the situation and distinguish between the dierent models, measurements of the ux of individual elements, or at least groups of elements, up to the highest energies are necessary. The range of cosmic rays ux is so wide that it requires several dierent detection techniques over the whole energy spectrum. At low energies the ux is large enough that sophisticated detectors with an active area of a few 0 cm 2 can be built to measure the abundance of individual isotopes. An energy of 14 ev is about the upper limit above which it is no longer possible to detect cosmic rays directly, with detectors in space or at the top of the atmosphere. Above this energy, in fact, the ux becomes so low to require detectors with large areas exposed for long periods of time. The only way to satisfy these requirements, at present, is to build

27 Cosmic rays 23 extended groundbased experiments to measure secondary products generated by cosmic rays in the atmosphere. Studying the cascade of particles, called Extensive Air Shower (EAS), it is possible to obtain indirect information about the nature of primary particles. In this way the (average) mass can be estimated only with large uncertainties. In fact, the interactions in the atmosphere destroy much of the information about the nature of the primary particle and extensive simulations are required to identify the composition of the primary particle ux. The situation is sketched in Figure 1.. Composition resolution isotopes (Z,A) Si detectors elements + - Z B spectrometers air Cerenkov elements Z calorimeters, TRDs emulsions elemental groups ln A air shower arrays MeV GeV TeV PeV EeV Energy Flux dn/de [m-2 sr-1 s-1 GeV-1] Figure 1.: Illustrative sketch of the composition resolution achieved by dierent cosmicray techniques as a function of energy. Over the energy range shown, the ux of cosmic rays decreases by about 30 orders of magnitude as indicated on the righthand scale Indirect measurements To have access to energies above 14 ev 15 ev, very large exposures are necessary. At present this task is reached only in groundbased experiments, recording extensive air showers. Extensive Air Showers A high energy primary particle enters the atmosphere and interacts with air molecules initiating a cascading process that produces secondary particles. A cosmicray induced shower in the atmosphere is made up of three components, i.e. electromagnetic, muonic and hadronic (see Figure 1.11(a)). The shower consists of a core of high energy hadrons, continuosly feeding the electromagnetic component of the shower, mostly by neutral pions and by eta particle decays into photons. An electromagnetic cascade is generated from each high energy photon by a long alternate sequence of pair production and bremsstrahlung, starting at its injection point. The muonic component is fed up by decay of lower energy charged pions and kaons. The electromagnetic cascade is the dominant component in the shower and dissipates most of the primary energy. In fact, at each

28 altitude (' 40 km) 24 Cosmic rays primary CR atmospheric nucleus e + e? e + 0 e?? + (?),K, N,: : : hadronic shower +? e? e e? electromagnetic shower e + Cherenkov radiation sea level e (a) (b) Figure 1.11: Schematic view of an extensive air shower. (a) Development of the dierent components; (b) spatial development in the atmosphere. hadronic interaction, slightly more than one third of the energy goes into the electromagnetic component and most hadrons undergo more than one interactions. Moreover, the increase of the electromagnetic cascade is so fast that electrons and positrons soon become the most numerous particles in the shower. After reaching a maximum, the number of electrons and positrons quickly decline, because, below the critical energy in air, E c 80 MeV, electrons lose their remaining energy into ionization. Conversely, the number of muons increases, as the shower develops, nally reaching a plateau, because muons do not interact but only lose energy by ionization as they traverse the atmosphere. Detection principles In the TeV regime (small) air showers are observed with imaging Cherenkov telescopes. These instruments image the trajectory of a shower in the sky with large mirrors onto a segmented camera. When it reaches the ground, the footprint of an air shower can cover an area of tens of square kilometers (see Figure 1.11(b)). The secondary particles also collide with and excite nitrogen molecules in the air, and thereby provide a ash of uorescence light 1 in the atmosphere. In today's experiments, the energy is basically derived from the density of secondary particles observed and the primary's mass is estimated by measurements of the depth of the shower maximum (measuring the longitudinal development of showers) or the electrontomuon ratio. There are three main types of detectors: the ground array, the air uorescence detectors and Cherenkov telescopes. Ground array experiments sample the charged secondary particles as they reach the 1 the light being emitted by the deexcitation of the nitrogen molecules

29 Cosmic rays 25 ground. The total number of particles at the observation level is obtained through the measurement of particles' densities and the integration of the lateral density distribution. In this way, the energy of primary particles is reconstructed. The direction of air showers is obtained through the measurement of the arrival time of the shower particles in the individual detectors (i.e. through trigger time). The primary chemical composition can be inferred from the electrontomuon component ratio in showers. A model of hadronic showers, due to Heitler, yields the relation: or N ( e E N µ A 1 PeV )0.15 (1.39) ( Ne ) log = C ln A (1.40) N µ where C is a parameter weakly dependent on the primary particle energy E 0. Another type of detector, the air uorescence detector, views tracks of light in the atmosphere. It determines the track geometry either via the photomultiplier tube trigger times or by the socalled stereo reconstruction, then calculates the primary energy by an integral along the track length and deduces the chemical composition by the shape of the longitudinal shower development. These experiments rely on the fact that the depth of the shower maximum for a primary particle of mass A relates to the depth of the maximum for proton induced showers as where χ 0 = 36.7 g/cm 2 is the radiation length in air. χ A max = χ p max χ 0 ln A (1.41) There is a copious Cherenkov light produced along the shower axis by the charged particles, and a large area Cherenkov array can be used to detect that light. The total ux of the Cherenkov light is a good measure of the total particle track integral in space and is thus a good primary energy parameter. The angular and lateral distribution of the Cherenkov light can be used to deduce the primary composition. To illustrate the sensitivity of air shower experiments to ln A one should bear in mind that to measure the composition with a resolution of 1 unit in ln A the shower maximum has to be measured to an accuracy of about 37 g/cm 2 or the N e /N µ ratio has to be determined with a relative error around 16%. Due to the large intrinsic uctuations in air showers, with existing experiments, at the most groups of elements can be reconstructed with ln A Experimental Results: an overview The rst experiment to claim the observation of a cosmicray particle above 20 ev was Volcano Ranch [25], in 1962 in New Mexico. It was the rst large shower array, made of 19 plastic scintillation counters, covering a total surface of 8.1 km 2. Born from the pioneering work of Linsley, Scarsi and Rossi, in 1959, it stayed operative until A few years later, a new large array detector [27] was built in Haverah Park, UK. It

30 26 Cosmic rays consisted of 32 water Cherenkov detectors for a total collecting area of 12 km 2. The Haverah Park detector took data from 1968 to SUGAR [26] has been, for a long time, the only giant array to be ever operative in the Southern Hemisphere. The Sydney University Giant Array Airshower Recorder was built by the University of Sydney at Narribri, New South Wales, Australia, close to the sea level, in It consisted of 47 detectors with a total collecting area of 70 km 2. The Yakutsk [28] array began taking data in 1969, in Siberia, and was developed to its full area of 18 km 2 four years later. In 1995 the array was contracted to km 2. The array utilized plastic scintillators, underground muon detectors and atmospheric Cherenkov light detectors to determine primary particle energy and composition. The rst extensive air shower detector based on uorescence technique was Fly's Eye [ 29], located in Utah, USA, about 160 km southwest of Salt Lake City. It was operated from 1981 to There were two uorescence detectors, Fly's Eye I and Fly's Eye II, able to perform either a monocular reconstruction, if the event was observed by only one of the two eyes, or the stereo reconstruction, if the event was recorded by both detectors. The integrated monocular exposure was about seven times larger than the stereo exposure, but the measurememnt accuracy greatly increased with stereo analysis. Fly's Eye has been replaced later by its updated version, HiRes [31], the High Resolution Fly's Eye. Until November 2003, the largest array constructed has been the Akeno Giant AirShower Array [30], AGASA, covering an area of about 0 km 2, at Akeno, Japan. It detects the cosmicray induced showers of charged secondary particles using a large network of surface plastic scintillators (111 detectors, separated about 1 km from each other), underground muon and water Cherenkov detectors. The KASCADE [32] (KArlsruhe Shower Core and Array DEtector) experiment is located at the site of the Forschungszentrum Karlsruhe at an altitude of 1 m above sea level. Its conguration is based on the concept of measuring as much as possible redundant information from each single airshower event. In fact, this multidetector system allows to measure the electromagnetic and muonic components of extensive air showers separately using an array of 252 scintillation detector stations, equally spaced by 13 m and covering an area of m 2. Additionally muon densities at further three muon energy thresholds and the hadronic core of the shower by a 300 m 2 iron sampling calorimeter are measured. The AUGER [33] project, named after Pierre Auger, the discoverer of extensive air showers, became operative only recently. Located in the Mendoza province in Argentina and with an active area of 3000 km 2, AUGER is the world's largest cosmicray detector. With a combination of water Cherenkov tanks and uorescence detectors, it is hoped that the AUGER project will put an end to the controversy surrounding ultrahigh energy cosmic rays. The above review of groundbased experiments is not exhaustive and only a few experiments are highlighted.

31 Cosmic rays 27 Experiment Types of detectors Exposure [km 2 sr year] No. of events Expected No. Area [km 2 ] Volcano Ranch scintillators muon detectors > 220 MeV SUGAR muon detectors > 0.75 GeV Haverah Park water Cherenkov Yatutsk scintillators up to 1995 muon detectors > 1 GeV Fly's Eye (mono) uorescence detectors Fly's Eye (stereo) uorescence detectors AGASA scintillators muon detectors > 1 GeV Table 1.1: Exposure, number of events and area for each experiment.

32 28 Cosmic rays Results from various experiments are compiled in Figure 1.12(a). At energies around 14 ev the values for the mean logarithmic mass are complemented by the results obtained from direct measurements. An increase of ln A as a function of energy can be recognized. However, individual experiments exhibit systematic dierences of about ±1 unit. Such uctuations in ln A are expected according to the simple estimate (1.39), assuming that the ratio of the electromagnetic and muonic shower components can be measured with an accuracy of the order of 16%. This uncertainty is a realistic value for the resolution of air shower arrays. In fact, the uncertainty in Monte Carlo models, describing the nucleusnucleus interactions during the air shower development, is reected in large systematic errors in the estimation of the cosmicray composition. Mean logarithmic mass <ln A> lg E 0 [GeV] BLANCA JACEE Fe G CASA MIA RUNJOB 3.5 Chacaltaya G DICE G G G G A Fly s Eye Mg 3 Haverah Park G d HEGRA (Airobicc) G d N HEGRA (CRT) A AA A A A 5G d d 5 5 HiRes Be 2 5 AA A G 5 d d d d G He G G G 1 d KASCADE (nn) 0.5 KASCADE (hadrons) KASCADE (electrons) SPACE 0 H Energy E 0 [GeV] X max (g/cm 2 ) 00 DICE [17] HEGRA [18] 900 CASA-BLANCA [19] CACTI [55] SPASE-VULCAN [56] Yakutsk [57] 800 Fly s Eye [58] HiRes-MIA [59] Haverah Park [60] SIBYLL 2.1 [61] 700 QGSJET [62] log ( E/eV ) (a) (b) Figure 1.12: (a) Mean logarithm of the primary particle mass as a function of the energy for selected cosmicray experiments. (b) Average depth of the shower maximum χ max as a function of primary energy for selected cosmicray experiments. Another technique to determine the mass of cosmic rays is to measure the average depth of the shower maximum χ max. The results of several experiments are presented in Figure 1.12(b). It depicts the measured χ max values as a function of energy; the upper pair of lines indicates showers initiated by protons and the lower pair is for ironinitiated showers. It is evident that there is a trend for a transition from heavier components to a larger fraction of protons around 15 ev, followed by a transition back toward heavier elements above the knee. The data in Figure 1.12(b) show systematic dierences of 30 g/cm 2 around 1 PeV increasing to 65 g/cm 2 close to PeV. In this case too, some of the experimental uncertainties may be caused by application of dierent simulation codes for air shower development in the atmosphere Direct measurements The detection of cosmic rays above the atmosphere is the only way to obtain direct measurements of the primary particles and their energy spectra. There are a number of techniques for direct detection. At energies in the MeV range, Lithiumdrifted silicon detectors, operated in outer

33 Cosmic rays 29 space, can identify individual isotopes, fully characterized by simultaneous measurements of their energy, charge and mass. Since the particles have to be absorbed completely in a silicon detector, this technique works up to energies of a few GeV only. In the GeV domain, particles are registered with magnetic spectrometers on stratospheric balloons. Magnet spectrometers are the only detectors suitable to distinguish between matter and antimatter. The particle momentum is derived from the curvature of the trajectory in the magnetic eld, which limits the use of these detectors to energies approaching the TeV scale. At higher energies, particles are measured with balloon borne instruments on circumpolar long duration ights. Individual elements are identied, characterized by their charge and energy. In general, Z is determined through specic ionization loss de/dx measurements. Experimentally, the energy measurement is the most challenging: in this case, we can dene active and passive techniques. The socalled active techniques usually involve the use of scintillators combined with a calorimeter or transition radiation detector (TRD). TRDs rely on the passage of a particle through the detector without a nuclear interaction (the Lorentz factor γ of the primary particle is measured, a quantity which is related to the kinetic energy per nucleon), while calorimeters not only require an interaction but must also contain as much of the ensuing cascade as possible for accurate energy measurements. Due to weight limitations, actual calorimeters have to nd an optimum between detector aperture and energy resolution, resulting in relatively thin detectors. TRDs can be constructed from lighter materials than those used for calorimeters. Hence, TRDs can have greater geometrical factors for a given weight compared with calorimeters. However, they suer from limitations due to the poor light yield and saturation eects. The former causes their detection eciency not being sucient for nuclei with Z < 3: only calorimetric techniques can be used to obtain reasonable energy measurements of protons and helium up to knee energies. The latter sets an upper limit to the maximum particle energy that can be measured with a given radiator and detector geometry. Current advances in TRD design [48, 49, 50]have shown that their operational range can be extended to energy measurements with Lorentz factor 4 γ 5. Nuclear emulsion is an example of a passive technique that, when interleaved with suitable converters, can simultaneously measure the primary charge and energy. These payloads must be recovered and subjected to complicated procedures to extract the energy and charge information. While each technique has its own strengths and weaknesses, the major limitation to all direct techniques to date is the collecting power. Due to the limited uxes at high energy, largeacceptance detectors are required to collect a suciently large data sample within an acceptable length of time. This limit is constantly being pushed as with the incoming advent of a new generation of balloons, the Ultra Long Duration Balloons (ULDB), designed for 900 day ights, and by constructing payloads that do not require pressurized vessels for the operation.

34 30 Cosmic rays Spacecraft instruments Spacecraft experiments are able to perform energy and charge measurements free of the overburden of the atmosphere. Below is a brief description of some recent instruments. High Energy Astronomical ObservatoryC2 experiment (HEAO3C2) The FrenchDanish Cosmic Ray experiment C2 [40] was launched in 1979 onboard the NASA HEAO3 satellite. The HEAO3 mission performed a sky survey of gamma rays and cosmic rays. The scientic objectives of the experiment HEAO3C2 were the determination of the isotopic composition of the most abundant components of the cosmicray ux with atomic mass between 7 and 56, the measurement of the ux of each element with atomic number 4 Z 50 and the search for superheavy nuclei up to Z = 120. The normal operating mode was a continuous celestial scan spinning around an axis pointed towards the sun. The altitude was 500 km, the inclination of the orbit was 43.6 and the angular acceptance of the detector was a cone of 28 from the axis. The detector consisted of ve Cherenkov counters with dierent refractive indices and an hodoscope of four ash tube array inserted between the counters. The instrument being bidirectional, a time of ight measurement between top and bottom counters was used to identify the direction of propagation of the particle. Three inner Cherenkov detectors were used primarily for velocity determination, while the top and bottom counters for charge determination. The total geometric acceptance was 0.07 m 2 sr. HEAO3C2 made detailed isotopic composition of cosmic rays above 2 GeV/n, and measurements of the elemental composition and energy spectra from 4 Z 28 in the energy range from 620 MeV/n to 35 GeV/n. Cosmic Ray Nuclei experiment (CRN) The Cosmic Ray Nuclei experiment [41] was launched on the Space Shuttle in The CRN instrument employed the rst transition radiation detector (TRD) designed to measure the energy of cosmicray nuclei. Using gas Cherenkov counters and TRD, CRN measured the spectra of individual elements from boron to iron above 50 GeV/n up to energies of about 2 TeV/n. Because of limitations in the dynamic range of the detector and due to the inherent magnitude of uctuations in the TRD and Cherenkov signals, the CRN instrument was not designed to provide observations of protons and helium nuclei. Thus, the major primary nuclei observed are carbon, oxygen, neon, magnesium, silicon and iron. CRN obtained relative abundances of the elements that are quite similar to those reported at lower energy. In particular, the depletion of cosmicray source abundances with increasing rst ionization potential, relative to the local galactic abundances, seems to persist well up to TeV energies. The intensity ratios of the secondary to primary nuclei B/C and N/O and that of the primary species C/O are shown in Figure 1.13(c) as lled data points. Carbon and oxygen have very similar energy spectra, with spectral indices of 2.60 ± 0.1 and 2.65 ± 0.1 respectively. Consequently, the abundance ratio C/O slightly changes within the experimental uncertainties up to energies in the TeV/amu region. This behaviour is expected if both carbon and oxygen are essentially primary nuclei and are produced in the same sources. The spectra of boron and nitrogen, however, are much steeper;

35 Cosmic rays 31 therefore, the abundance ratio B/C and N/O continue to decrease up to energies around 1 TeV/amu: as discussed previously, this directly reects the mean escape path length of material traversed during the propagation of the particles. Alpha Magnetic Spectrometer (AMS) The Alpha Magnetic Spectrometer [42] is a high energy physics experiment scheduled for installation on the International Space Station. In preparation for this long duration mission, AMS had a precursor ight (AMS01) on board of the space shuttle Discovery (ight STS91), lasting days, from June 2 to 12, AMS01 has been the rst magnetic spectrometer in space with large collecting area ( 1 m 2 ) and with an analyzing power of 0.15 Tm 2. This permanent magnet contains 4 of the 6 layers of the siliconbased tracker and the scintillator counters of the anticoincidence system. Each tracker plane gives a measure of 2 coordinates (x,y) and of energy deposition, thus giving the particle momentum and charge. The trigger is given by the time of ight system, that in addition measures the velocity of the traversing particles and their charge. When combined to the tracker measurements, this allows the determination of the particle mass. The detector is completed with a threshold Cherenkov counter, below the magnet, to improve the separation between electrons and protons. The altitude is 380 km, the inclination of the orbit is 51.6 and the angular acceptance of the detector is a cone of 32 from the axis. AMS01 has provided detailed measurements of charged particles ux outside the Earth's atmosphere, collecting 7 protons in the energy range GeV, 5 electrons in the energy range GeV and 6 helium nuclei in the energy range 0.10 GeV/n. Above the geomagnetic cuto the observed spectra are parameterized by a power law in rigidity. The t yields γ H = 2.79 ± 0.01/ ± and γ He = ± 0.0/ ± Payload for Antimatter Matter Exploration and Lightnuclei Astrophysics (PAMELA) The PAMELA [43] experiment is a satelliteborne apparatus designed to study charged particles in the cosmic radiation with a particular focus on light nuclei and antinuclei in the energy range from 8 to 11 ev/n, as well as electrons and positrons up to ev. The apparatus comprises a timeofight system, a magnetic spectrometer, a silicontungsten electromagnetic calorimeter, an anticoincidence system, a shower tail catcher scintillator and a neutron detector. The combination of these devices is expected to allow antiparticles to be reliably identied from a large background of other charged particles. PAMELA has been launched on a Russian satellite in June 2006 and is taking data successfully Current and recent balloon instruments Since satellite experiments are often quite expensive and limited to smaller and lighter detectors, balloonborne instruments are reliable alternatives that can achieve similar scientic goals at a fraction of the cost. While the exposure time for balloon experiments is limited to tens of days, balloonborne detectors can be larger and heavier with greater

36 32 Cosmic rays Experiment Description Period Cherenkov - hodoscope reference set for HEAO3C2 4 Z 28 secondarytoprimary ratio analysis Cherenkov - TRD 1985 few hours ight time CRN 5 Z 28 rst TRD in space Magnetic Spectrometer 1998 days ight time AMS01 calorimeter rst magnetic spectrometer in space Table 1.2: knee. Selected spacecraft instruments which measured cosmicray composition below the geometric factors than those achieved in spacebased experiments. The ideal locations for balloonborne cosmicray detectors are at high latitudes, near the north and south poles, where the geomagnetic cuto is low and balloons can remain aloft for several weeks without invading commercial airline space or oating above populated areas. Finally, besides securing scientic data, balloon experiments can be used as test prototypes for future satellite experiments. JACEE [34] and RUNJOB [35] were passive emulsion/xray lm calorimeters which ew 15 and 11 times respectively in order to accumulate the needed high energy statistics. Passive detectors have the advantage that the equipment is relatively inexpensive, and the detectors can be duplicated and own multiple times in order to build up a large exposure factor. Both JACEE and RUNJOB experienced an unsuccessful campaign. Nevertheless, by combining the results from multiple campaigns, JACEE and RUNJOB have provided the highest energy direct measurements of the hydrogeniron spectra that are currently available. JACEE has reported results based on a sample of 656 proton events above 6 TeV and 414 helium nuclei above 2 TeV/nucleon, and RUNJOB has detected a proton event with energy near 00 TeV. In the following, we shortly summarize the main features of a selection of recent balloon experiments using the active detection techniques (Table 1.3). The important series of balloon experiments from the BESS collaboration, mainly dedicated to antiproton measurements and antimatter search, is not included in the following. ATIC [36] is a large (0.24 m 2 sr) calorimeter with signicantly greater depth than the emulsion instruments (22 radiation lengths X 0 compared to 9X 0 for JACEE). TRACER [38] is the largest balloon instrument: it is a m 3 transition radiation detector. CREAM [37] is the rst instrument which combines a transition radiation detector with a calorimeter to provide independent energy measurements of cosmicray nuclei. TIGER [39] is an instrument designed to detect nuclei well above iron. Though TIGER not providing data at extremely high energies, it is included here because a complete picture of the cosmic ray composition requires extending current measurements both in energy and charge.

37 Cosmic rays 33 Experiment Description No. balloon Charge Range Aperture [m 2 sr] ights Series of emulsion 15 (12) 644 m g/cm 2 JACEE experiments Z Zenith angle acceptance m 2 sr days exposure Series of emulsion 11 () 575 m g/cm 2 RUNJOB experiments Z Highest energy proton event E > 1 PeV Silicon matrix - 3 (2) 31 days ight time ATIC scintillator Z m 2 sr days exposure BGO calorimeter Scintillator Z days ight time TRACER Cherenkov - TRD m 2 sr days exposure 5 Z 26(2006) Scintillator - TRD Z 28 CAL days ight time CREAM Si charge detector Goal is to y multiple ULDB W+scintillator CAL 4 Z 28 TRD 1.3 ights to build up exposure Scintillator m 2 sr days exposure TIGER Cherenkov - ber Z Originally planned for ULDB hodoscope ight: future ights planned Table 1.3: Current and recent balloon instruments measuring high energy cosmicray composition. The experimental exposure is calculated for the eective number of ights (indicated in brackets) used for the analysis.

38 34 Cosmic rays Japanese American Cooperative Emulsion Experiment (JACEE) The JACEE collaboration ew a series of thin emulsion/xray lm calorimeters on 15 ights, including 5 longduration Antarctic ights (JACEE 14). Each emulsion chamber was composed of four sections: the primary charge detector, the target, the spacer and the calorimeter. The primary charge detector contained both thick and thin nuclear emulsions, which allowed accurate determination of the charge of each primary particle. The target section employed about fty thin emulsion plates interleaved with low Z material to maximize the nuclear interaction probability, while minimizing the probability for pair production by high energy photons. The spacer section was a drift space used to facilitate the tracing process and to allow measurement of charged particle emission angles as small as 5 rad. The calorimeter section had about 79 radiation length vertical thickness of lead absorber. Each lead sheet was interleaved with Xray lms and thin emulsion plates. Its role was to collect the fraction of incident particle energy which appeared as γrays in the interactions. The average ight altitude ranged from 3.5 to 5.5 g/cm 2. The total number of high energy events available for analysis from all ights was 2 4. Of these, 180 had energy exceeding 0 TeV/particle. JACEE's published results on the hydrogen and helium spectra were based on data from 12 of those ights (JACEE 11 was lost in the sea, and the data from JACEE were never analyzed). The binned data extended earlier results to energies close to 200 TeV for protons and 0 TeV/nucleon for helium nuclei (see Figure 1.13(a)). In order to eliminate the uncertainty due to the arbitrary binning of data with small numbers of high energy events, the JACEE group presented integral spectra (see Figure 1.13(b)). A likelihood analysis gave best t integral spectral indices for the protons and helium that diered at the 2σ level: γ H = 2.80 ± 0.04 and γ He = 2.68 ± 0.04/ Although the likelihood analysis was designed to eliminate correlations, the waviness observed in the integral spectra in Figure 1.13(b) could be an indication of systematic errors resulting from combining data from multiple ights with slightly dierent normalizations and calibrations. Actually, JACEE remains the balloon experiment with the highest statistics at energies above TeV. RUssianNippon JOint Balloon collaboration (RUNJOB) RUNJOB ew a more or less similar set of Xray lms and emulsion chambers on a series of 11 balloon ights from Kamchatka starting in 1995, of which were successful with a total of 575 m 2 sr of exposure. Although RUNJOB ew somewhat lower in the atmosphere than JACEE ( g/cm 2 ), RUNJOB has presented results for hydrogen through iron. The integral spectral indices for the RUNJOB hydrogen and helium are equal to within the statistical errors: γ H = 2.74 ± and γ He = 2.78 ± 0.20 at energies up to 0 TeV. In the 1995 campaign, RUNJOB detected a single proton with energy near 00 TeV and, over a period of ights, saw no evidence for any steepening in the proton spectrum with increasing energy. The absolute intensity of the hydrogen measured by RUNJOB agreed reasonably well with the intensity measured by JACEE, but RUNJOB reported a helium ux which is lower than that of most other experiments by about 40% (it is in agreement with that reported by MUBEE). For charges heavier than helium, RUNJOB presented the energy spectra of three charge

39 Cosmic rays 35 groups: CNO, NeMgSi, and Fe, with a slope reported to harden from 2.7 for CNO to 2.6 for Fe (see Figure 1.13(d)). The resulting average mass ln A reported by RUNJOB was roughly constant at 2.5 up to 00 TeV/nucleus, in reasonable agreement with the results from the KASCADE airshower experiment and the BLANCA Cherenkov light measurement, but in disagreement with the results from CASAMIA, which showed ln A increasing with energy starting well below 00 TeV/nucleus. Advanced Thin Ionization Calorimeter (ATIC) ATIC is an active instrument combining: a negrained silicon matrix for charge measurements a 0.75 interaction length graphite target to induce nuclear interactions scintillator strip hodoscopes for triggering and trajectory reconstruction an 18 radiation length deep (22X 0 for ATIC3) bismuthgermanate (BGO) calorimeter to measure the energy of incident particles, with a geometry factor of 0.24 m 2 sr. ATIC is intended to provide sensitivity for hydrogen through nickel from 50 GeV to 0 TeV total energy. ATIC has own three times from the Antarctic with a total 31 days exposure time; ATIC, too, experienced an unsuccessful campaign in December Measurements from JACEE suggest a steeper spectrum for proton with respect to helium. This dierence is not conrmed by RUNJOB and ATIC1 data which are instead consistent with a similar spectral index for the two lighest elements. However, the preliminary ATIC2 results for both proton and helium nuclei show again (see Figure 1.13(e)) dierent spectral indices γ H 2.9 and γ He There is a reasonable agreement with the data of magnetic spectrometers in the lowenergy region. The dierences between the results of the two ights appear to be due to a better statistics and an improved procedure for tting the ATIC2 spectra. In the case of the heavier nuclei, ATIC has sucient charge resolution to identify individual charges. This is in contrast to the emulsion instruments, which showed nuclear results for only charge groups. The ATIC2 nuclear spectral results are shown in Figure 1.13(f), where the statistics are clearly not yet sucient to be able to make strong claims at energies above 1 TeV/nucleon. Transition Radiation Array for Cosmic Energetic Radiation (TRACER) TRACER is the largest of the current generation of balloonborne detectors. It utilizes a m 3 high array of transition radiation detectors together with a set of plastic scintillators and an acrylic Cherenkov counter to measure the charge and energy of nuclei with Z 8 during the rst ight and Z 5 during the second ight. The relatively low density of the transition radiation detector with plastic ber radiators and multiwire proportional chambers makes it possible to design the instrument with a very large collecting power ( 5 m 2 sr) needed to perform measurement of the energy spectra for

40 36 Cosmic rays individual elements at energies up to 14 ev. TRACER has own three times: a 30 hour test ight from Ft. Summer, a 14 day exposure over Antarctica (December 2003) and a 4.5 days journey around the Artic Circle from Sweden to Alaska (Summer 2006). The lastest ight is marked by the extension of the instrument's sampling range to encompass elements from boron to iron. Absolute intensities and the Fe, Si, and Mg spectral slopes measured on the initial ight agree well with previous observations in space with CRN and HEAO, although the preliminary Ne and O spectra may be slightly atter than those reported from CRN (see Figure 1.13(f)). Cosmic Ray Energetics And Mass (CREAM) CREAM combines silicon charge detectors, scintillating ber hodoscopes, a cm 3 transition radiation detector, and a 20 radiation length cm 2 tungsten scintillator sampling calorimeter. The calorimeter provides sensitivity for hydrogen and heavier particles, and the transition radiation detector is designed to be used for Be to Ni. The instrument has been designed for 0day Ultra Long Duration Balloon (ULDB) ights, with a goal of ying multiple times in order to obtain statistically signicant spectral and composition information for protons through iron from 12 to 15 ev. CREAM was successfully own twice on longduration balloons from McMurdo, Antarctica in 2004/05 and 2005/06 with about 70 days total exposure at an altitude of 3540 km. For a detailed description, see next chapter. TransIron Galactic Element Recorder (TIGER) TIGER was designed to probe high Z cosmic rays. The instrument is approximately a 1 m 2 detector which uses a combination of scintillators together with plastic and Aerogel Cherenkov detectors and a scintillating ber hodoscope to measure the charge and energy of incident nuclei for charges from Z 26 to 40 at energies from 800 MeV/n up to approximately GeV/n. TIGER was the rst payload to be selected for ULDB ights, but because of delays in the ULDB program, it was own three times on shorter ights, once for one day from Ft. Sumner (1999) and twice from the Antarctic (2001 and 2003) for a total of 50 days of additional exposure. Measured abundances appear to be in good agreement with solar system abundances modied by either rst ionization potential or volatility, but low statistics prohibits conclusive evidence of either a FIP or volatility interpretation of the data for galactic cosmic rays.

41 Cosmic rays 37 (a) (b) (c) (d) 1.75 (GeV) -1, m -2 s -1 ster 2.75 Flux E 5 ATIC, p 4 AMS, p CAPRICE98, p BESS-TeV, p ATIC, He AMS, He CAPRICE98, He BESS-TeV, He 1.75 (GeV) -1, m -2 s -1 ster 2.75 Flux E ATIC HEAO-3-C2 CRN TRACER C O 0.01 Ne ATIC HEAO-3-C2 CRN TRACER Mg Si 0.01 Fe Energy per particle, GeV Energy per particle, GeV Energy per particle, GeV 5 6 (e) (f) Figure 1.13: (a) JACEE dierential spectra for proton and Helium. (b) JACEE integral spectra for proton and Helium. (c) The ratios B/C, N/O and C/O as a function of kinetic energy. Compilation of spectral data for (d) element groups CNO, NeMgSi and Fe; (e) for proton and helium (f) for separate evenz most abundant elements.

42

43 Chapter 2 The CREAM Detector The present chapter decribes the CREAM experiment. CREAM is a multimission program, which has had two successful ights in 2004/05 and 2005/06. Though this work reports preliminary results from the second ight, each subdetector, from the two congurations, will be described. Particular emphasis will be devoted to the calorimeter (which was built by the INFN group), since, for the second ight, it is the only instrument providing energy measurements. 2.1 Science objectives The Cosmic Ray Energetics And Mass (CREAM) experiment is designed to perform direct measurements of cosmicray composition and energy spectra above the atmosphere. The goal of this multimission experiment is to collect, in a series of ights, sucient statistics to explore the region of energies from 12 up to 15 ev over the elemental range from hydrogen to iron (1 Z 26), thereby providing calibration for the simulation models used by groundbased airshower experiments near the low end of their energy range. One of the main physics objectives driving the CREAM investigation is the experimental test of the validity of models based on cosmicray acceleration by supernova shock waves. In order to achieve this goal, CREAM must search for a possible break or bend in the spectra of light nuclei; search for the consequent change in the elemental composition of cosmic rays, as a function of energy; measure the energy dependence of the protontohelium ratio; assess possible dierences in the spectral slopes of heavier nuclei with respect to helium and to protons. A common feature of all the models is their capability to predict the detailed shape of the energy spectrum of each element [52]. Therefore, in order to discriminate among dierent models, CREAM has to provide an unambiguous identication of the incoming particle (via a precise measurement of its charge) and a suciently accurate determination of its energy. Another broad science objective, addressed by CREAM, is the measurement of the spectra of secondary elements up to the highest possible energies allowed by the irreducible background, generated by the residual atmospheric overburden at balloon altitudes. The 39

44 40 The CREAM Detector aim is to get a consistent picture of the propagation of cosmic rays through the Galaxy and of their interactions with the interstellar medium. Measurements of secondaryto primary ratios [54] would allow the determination of the parameter δ whose value is critical to discriminate among dierent models [53, 55] predicting a possible E δ energy dependence for the propagation pathlength. Once the value of this parameter is known with sucient accuracy, it will be possible to infer the shape of the acceleration spectra of individual elements at the source. These science objectives drive the following measurements goals: the CREAM calorimeter is designed to collect as many high energy proton and helium events as possible; CREAM data analysis is expected to reconstruct 100 TeV primary energies; the energy resolution should be better than 50% at all energies of interest; CREAM analysis is expected to reconstruct the primary particle's charge well enough to identify individual elements; CREAM is designed with a geometry factor sucient to collect Boron and Carbon data with sucient statistics to reconstruct the B/C ratio up to 500 GeV and to accumulate data on CNO, NeSi and Fe group nuclei above 14 ev. 2.2 Instrument design Originally planned to be own on the new Ultra Long Duration Balloon (ULDB) being developed by the NASA, the CREAM instrument was instead launched twice on a long duration balloon (LDB) from McMurdo Station, Antartica. Therefore, in order to maximize the total exposure of the experiment as a multimission program, two instrument suites were built to be own on alternate years (December 2004 and December 2005). In this way, the refurbishment operations that follow the recovery of the rst payload could take place almost simultaneously with the ight preparation of the second payload. The CREAM instrument must determine the charge Z and energy E of the incoming particles to meet the science objectives. For this purpose, as shown in Figure 2.1(a), it is comprised of several major detector subsystems. In the CREAMI conguration, four independent charge measurements of incident particles are performed using (from top to bottom) a Timingbased scintillator Charge Detector (TCD), a plastic Cherenkov Detector (CD), a Silicon Charge Detector (SCD) and scintillating ber hodoscopes (S0 /S1 and S2). In fact, one of the consequences of having a high density calorimeter (for energy measurement) as part of the experiment is the large number of secondary shower particles generated in the calorimeter absorber, some of which are scattered back up towards the top of the payload. To overcome this challenge, the CREAM design includes redundant charge measurement systems, accomplished principally either by detector segmentation or by a pulseshape timing technique. The TCD utilizes a fast readout system that takes advantage of the brief time interval between the primary's traverse and the arrival of the rst secondaries to complete its charge measurement, thereby preventing the back scatter noise impacting its measurement of primary charge. The SCD is nely segmented

45 The CREAM Detector 41 TCD TRD Modules SCD Hodoscopes Calorimeter S3 detector Graphite Targets (a) (b) Figure 2.1: (a) Schematic cross sectional drawing of major CREAM detector systems. (b) CREAMII launch.

46 42 The CREAM Detector into pixels of 2.12 cm 2 area each, which minimizes the number of backscattered particles hitting the same pixel that measures the primary particle. The CD responds only to (relativistic) particles with velocity exceeding the velocity of light in the plastic, so it allows rejection of the abundant low energy cosmic rays present in the high latitude regions where the geomagnetic cuto is low. The ber hodoscopes also use spatial segmentation, by reading out separate bers, and provide additional charge measurements as well as particle tracking information. The particle energy is measured with a Transition Radiation Detector (TRD) and a tungsten/scintillator sampling calorimeter (CAL). For particles with Z 3, the TRD can measure the Lorentz factor γ. When paired with the knowledge of the particle mass, via charge identication, it provides an energy measurement independent from the calorimeter. The relatively low density of the TRD allows the larger geometric coverage needed for lower ux particles, with an acceptable detector weight. The calorimeter is preceded by a densied graphite target (T1 and T2) followed by a single ber layer shower detector (S3) read out by TCDtype photomultiplier tubes (PMT) to provide a reference time for the TCD. It can determine the particle energy for all nuclei including those with Z < 3 for which the TRD cannot reliably provide an energy measurement. The TRD and calorimeter have dierent systematic biases in determining the particle energy. The use of both instruments allows inight crosscalibration of the two techniques and thus a more reliable energy determination. While both detector types have been own separately in prior balloon investigations, CREAMI is the rst balloon payload to employ a calorimeter and TRD together for high energy composition measurements. The CREAMII instrument conguration does not include the TRD and scintillating ber hodoscopes as in the rst ight, while the charge identication performance of SCD are enhanced by the addition of a second layer of pixelated silicon sensors Timing Charge & S3 Detectors The TCD, which fully covers the top of the instrument, provides charge measurements for particles that enter the TRD acceptance. Accurate identication of the incident particle charge is hampered by albedo from the calorimeter, especially for lowercharged nuclei. Most charge detectors own on similar experiments have thus been designed with ne spatial segmentation (e.g. pixels or strips) that minimize the area over which measurements are integrated. The CREAM TCD [57] instead takes advantage of a unique property of the primary particle. For events with a shower in the calorimeter, about 75 cm below the TCD, one expects backscattered shower secondaries to traverse the charge detector several nanoseconds after the primary has gone through. This background noise is nearly proportional to the logarithm of the incident particle energy. For the geometry of the CREAM instrument, the arrival of the light signal from a primary nucleus at the TCD photomultiplier tubes precedes the earliest possible arrival of a light signal from a back scatter particle by 3 ns. Thus, using an ultrafast readout one can separate the primary signal from backscatter noise without resorting to high channel counts. Figure 2.2(a) shows a simulated 0 TeV proton event as seen in the time domain at each photomultiplier face in a TCD paddle. This gure illustrates both the signicant amount of albedo present and the clean separation of the initial particle in the time domain. The TCD is comprised of two crossed layers of four fast plastic scintillator paddles each,

47 The CREAM Detector 43 (a) (b) Figure 2.2: (a) Simulated energy deposits in a TCD scintillator due to a 0 TeV cosmic ray proton and the albedo generated by the subsequent shower in the calorimeter. The lower panels display the initial portion of the signal on an expanded timescale, showing the isolated 2 MeV pulse from the primary particle.(b) Schematic TCD paddle structure. read out with fast timing PMTs via twistedstrip adiabatic light guides. Each layer covers a 1.2 m 1.2 m square. Each Bicron BC-408 paddle is 1.2 m long, 0.3 m wide and 5 mm thick. The paddles in the two layers are oriented with their long axes perpendicular to each other. The light guides gradually change crosssectional shape from a narrow rectangle at the scintillator gluejoint to a nearly round crosssection where they are glued to the fast PMTs, with their crosssectional area remaining constant (see Figure 2.2(b)). The light guides are UVA acrylic, where the UV absorption reduces the eects of Cherenkov emission in the light guides without signicant reduction of the blue scintillation signal. The paddles are wrapped in crinkled aluminum foil and the wrapping is covered with DuPont black Tedlar to assure lighttightness. In order to determine the charge of incident cosmic rays with sucient accuracy ( 0.2e for O to 0.35e for Fe) to resolve individual elements, the amplitude of the scintillation signal is measured at four dierent dynodes (with correspondingly increasing dynamic range). The time structure (with 50 ps time resolution) of the leading edge of the light pulse at the anode is also measured. The timing readout is based on comparators, with multiple timetodigital converter (TDC) circuits digitizing the time at which the leading edge of the scintillation light pulse crosses two known thresholds. From the timing signals the TCD data processing code reconstruct the slope of the pulses leading edge, which is proportional to the square of the incident particle charge. The pulseheight readout captures the peak level and digitizes it using an analog todigital converter (ADC) circuit. This too is proportional to the square of the incident charge.

48 44 The CREAM Detector Scintillator Clamp Button (16 total) HV and IF Boxes TRD Amplex Electronics Box TCD Top Scintillator TCD Light Guide (16 total) TCD PMT and Electronics Large GAK Small GAK Cherenkov Detector Outrigger Support for TCD TRD Tray Figure 2.3: Schematic view of TCD, TRD and CD. The rst technique provides a charge measurement adequate to resolve the elements up to Z = 16, while the second readout is more accurate for higher charges where the primary particle is expected to have the highest peak value. The S3 detector provides a reference time for the TCD. Positioned directly over the calorimeter stack, it is a single layer of 2 2 mm 2 square, multiclad scintillating bers. The bers are all collected into a roughly circular array on either end, where a Lucite cylinder is glued onto the ber faces. The other ends of these light mixers are mated to the same type of PMTs and readouts as described above. The TDC digitizes the time at which the signal arrives on either end of S3. From this time the TCD code reconstructs the time at which the primary traversed the upper TCD layers. The TCD electronics is designed to complete the charge measurement within 23 ns, thereby avoiding the impact of background from backscattered particles, which arrive at the TCD 3 8 ns after the initial passage of the parent cosmic ray Transition Radiation & Cherenkov Detectors Immediately below the TCD is a Transition Radiation Detector (TRD). It measures the Lorentz factor γ of incident particles and can also reconstruct the particle trajectory through the detector. The CREAM TRD is comprised of two modules separated by a Cherenkov detector (CD, see description below). Each of the two TRD modules consists

49 The CREAM Detector 45 (a) (b) Figure 2.4: (a) Exploited view of a TRD thin walled proportional tube.(b) Schematic diagram of the Cherenkov detector with the plastic wavelength shifting bar readout. of four layers (in each transverse orientation) of polystyrene foam radiator combined with a total of m long thinwalled proportional tubes, each 2 cm in diameter (see Figure 2.4(a)). Each module is cm 3 in active volume. The foam, aside from providing mechanical support for the tubes, generates the transition radiation (TR) emitted by the charged primary particle as it repeatedly moves into and out of media with dierent dielectric constants. The tubes are lled with a mixture of xenon (95%) and methane (5%) gas at a pressure of 1 atmosphere to detect xray transition radiation in the energy region around kev. The tubes are constructed from aluminized Mylar wound with a wall thickness of 75 µm to allow easy penetration of transition radiation xrays. A central sensewire is installed in each tube, which is operated as a conventional proportional counter to detect ionization in the tube gas. The sensewire is strung between the two endcaps kept at 1.5 kv with respect to the grounded tube wall. The endcaps have been specially designed to minimize gas leaks to a level comparable to the leak rate through the Mylar walls. The TR signal generates charge pulses in the wires which are read out using an analog system based on the Amplex VLSI charge amplier chip, with a highgain and a lowgain range to cover the 12-bit dynamic range required to measure particles from Lithium to Nickel. The tubes are arranged in 8 layers (two in X orientation, two Y, two more X and two more Y for each module) and the pattern of hits in the tubes due to nuclei travelling through them is analyzed to reconstruct the particle's trajectory in three dimensions. A simple weighted line t to these hits provides a tracking (RMS) resolution better than σ 5 mm. A further t to the relative pulse heights in the tubes can improve this accuracy to better than 2 mm. The Lorentz factor γ of a heavy nucleus can be determined from the energy loss per unit pathlength in the TRD gas. In the region from minimum ionizing to 500 GeV/n, this is provided by the logarithmic relativistic increase in ionization loss, which is large ( 1.6 ) for xenon gas. At higher energies, above 1 TeV/n, signicant transition radiation is produced in the radiator material and accompanies the direct particle energy losses, providing an additional logarithmic rise in response until saturation sets in near γ 2 5. The Ethafoam radiator is optimized to cover the range 3 < γ < 5. The response of this TRD has been calibrated in a CERN test beam in 2001 [58]. In order to reject lowenergy particles for data collected near the South Pole (lower

50 46 The CREAM Detector geomagnetic cuto) a plastic Cherenkov layer has been added between the two TRD modules (see Figure 2.4(b)). The Cherenkov detector is a 1 cmthick 1.2 m 1.2 m plastic radiator sheet doped with blue wavelength shifter. This is viewed along the edges by eight small photomultiplier tubes via plastic wavelength shifting bars, which capture the blue light produced in the radiator and shift it into the green region. This technique provides a compact detector with uniform response. The measured variations of light collection with position in this device during the rst ight are < 15%. In this way the primary can be required to be relativistic, reducing the trigger rate for nonshowering events to an acceptable level < 30 Hz. The Cherenkov detector also provides a complementary determination of the incident cosmic ray charge with the TCD measurement. The heavy nucleus (HiZ) trigger of CREAM is set by thresholds in the TCD and Cherenkov detectors. During the ight these were adjusted so that a vertical relativistic Boron nucleus is well above these thresholds. The eective trigger aperture of the heavy nucleus trigger is fairly large with an acceptance of 2.2 m 2 sr, providing a large sample of secondary nuclei Silicon Charge Detector The main purpose for the SCD is to identify the charge of the incident cosmic ray particle (by measuring the rate of its ionization loss in the silicon) in the presence of backscattered particles from showers in the target and calorimeter below the charge detector. This is accomplished by a nely segmented array of pixelated PIN diode silicon sensor. With this segmentation, backscatter is expected to cause charge misidentication of only 2%3% of low-z particles at the highest incident energy ( 15 ev). For the CREAMII ight, the SCD has been upgraded with the addition of a second layer of sensors. The base material of the PIN diode silicon detector is a 380 µm thick, Ntype wafer with high resistivity (> 5 kω cm). Three guard rings have been implemented around active cells in the sensor design in order to reduce the electric eld on the side of the junction edge. The eect is conrmed in a reduced leakage current and improved junction (a) (b) Figure 2.5: (a) Assembly of sensors, analog board and thermal strap.(b) Fully assembled SCD for CREAMII ight.

51 The CREAM Detector 47 breakdown voltage. The fabrication procedure has been optimized for high yield of good quality sensors. A common contact is made on one side by phosphorous diusion while the individual detector pixels are created on the other side by boron ion implantation. The electrical characteristics like capacitance and reverse current at the various bias voltages have been checked for all sensors. For most of the sensors the full depletion voltage has been measured to be 80 ± 15 V and the leakage current has been measured to be 5 ± 3 na/cm 2 at 0 V and 25 C. The SCD [59] is comprised of 26 ladders mounted on an aluminum frame for the mechanical support. Each ladder holds 7 silicon sensor modules with associated analog readout electronics. The sensors are slightly tilted (1.24 degree from the horizontal line) and overlap each other in both lateral directions, providing full coverage in a single layer. Each silicon sensor is comprised of 16 silicon pixels in a 4 4 array, with 2.12 cm 2 active area (15.5 mm 13.7 mm) and a total of 2912 channels for the CREAMI single layer conguration. The total size is mm 3 and the coverage area of the SCD is mm 2. The total detector height is 35.2 mm including the detector cover. For CREAMII ight new sensors have been produced with nearly the same dimension ( cm 2 ). In this case, the postfabrication process has been improved for new sensors. In fact, a ground area is added in the space between pixel readout lines for stable operation with reduced readout noise. Figure 2.5(a) shows the assembly of sensors with spacers, analog board and thermal strap. A total of sensors are used in a layer with two extra sensors at two corners of the array. The top layer is identical in the structure and is installed perpendicularly on top of the bottom layer. In this way all the area is completely covered by sensors without insensitive region while the inner area of 0.52 m 2 is covered by an eective dual layer of sensors. Figure 2.5(b) shows fully assembled SCD for CREAMII ight. A total of 156 sensors is used for each layer, corresponding to 4992 silicon pixels in both layers. The height between the bottom of the frame and the top cover is 97.5 mm. The readout electronics are designed around a 16channel ASIC for each sensor, followed by 16bit ADCs, allowing ne charge resolution over a wide dynamic range covering from hydrogen up to nickel. Copper thermal straps are attached to the cover in order to keep the detector within its operational temperature range, conducting out heat from the frontend electronics Fiber Hodoscopes There are three hodoscopes in the CREAMI instrument. S0/S1 A pair of hodoscopes (S0/S1) is located at the top of the upper target and is comprised of four crossed layers of Bicron multiclad scintillating bers (see Figure 2.6(a)). Each ber has a square crosssection of 2 2 mm 2, with a thin layer of white extramural absorber (EMA) painted over the external cladding layer. Each ber is cut and polished on both ends, with vacuumdeposition of aluminum on one end forming a mirror to increase the eective light yield and reduce the dependence of signal strength on position of particle traversing the ber. Each ber is glued on the readout side to a short segment of a clear

52 48 The CREAM Detector ber also painted with white EMA. The clear bers are bent and routed into an aluminum cookie mated against a 73pixel hybrid photodiode (HPD). Alternate bers are read out from opposite ends for mechanical reasons. Each layer is comprised of 360 tightpacked bers, covering an active area of about cm 2. The lower hodoscope S1 is identical to S0 in all respects, except that the readout end of each scintillator ber ends at the other edge of the active area, where a clear ber with identical crosssection is glued to the scintillating ber. The clear bers are then populated in the same cookies as the S0 scintillating bers. In this way, alternate layers have orthogonal ber orientation relative to each other, with 2 layers reading the X orientation, and 2 others the Y orientation. 64 of the 73 pixels are read out using four 16channel ASICs, with 30 pixels reading out bers from the upper layer pair, 30 reading out bers from the lower layer pair and 4 reading out signals from light emitting diodes (LED). The LED signals are used in laboratory to optimize the pixeltober alignment and in ight as a light source for aliveness testing and calibration. S0/S1 is used to aid track reconstruction, improving the lever arm for track tting, and to provide supplemental charge identication. S2 The third hodoscope, S2 (see Figure 2.6(b)), is located between the two graphite target sections. This detector is constructed in a manner similar to S1, but with shorter scintillating bers. Since the targets are shaped as trapezoids which become narrower lower down, the active area of S2 is smaller, about cm 2. The bers in this detector are the same type and are read out in a similar manner. The S2 clear bers are routed, in groups of 64, to 12 HPDs (3 per side), with three clear ber bundles allowing LED light to be inserted into pixels of the S2 HPDs. S2 serves as another source of tracking information and helps to determine whether an interaction occured in the upper target. S2 was built by the INFN group and own on CREAM rst ight Target and Calorimeter Limitations on the calorimeter mass for balloonborne or spacebased payloads led to the development of the concept of Thin Ionization Calorimetry with the use of a lowz target preceding the calorimeter. An inelastic interaction of the primary nucleus in the target initiates a hadronic shower containing a narrow electromagnetic core, generated by the decay of neutral pions, which is imaged by a negraned calorimeter. Target The densied graphite target is 19 cm thick ( 0.46 interaction length λ int and 1 radiation length X 0 ). This graphite has a density of 2.1 g/cm 3 (compared to 1.7 g/cm 3 for typical industrial graphite). The target ares out in the shape of an upsidedown truncated pyramid with a slant angle of 30 (see Figure 2.7). It is divided into two elements cemented into composite cages and interleaved with S2. As mentioned above,

53 The CREAM Detector 49 (a) (b) Figure 2.6: (a) The S0/S1 ight module before integration. (b) The S2 ight module during the nal assembly phase in Italy. the target must force at least one hadronic interaction to provide an energy measurement. A thickness of 1λ int of either a low-z material (e.g. C or Be) or high-z material typically used in calorimeters (e.g. Fe, Pb, W) provides similar interaction probabilities. However, a low-z target with a smaller value in g/cm 2 requires less mass leaving a greater mass allowance for the calorimeter itself, allowing greater lateral size and therefore increasing the eective acceptance. Another advantage of low-z materials is the low ratio of λ int to X 0 which minimizes the dependence of calorimeter response on the rstinteraction depth. Calorimeter The calorimeter design is based around: a longitudinal segmentation in layers of 1 X 0 thickness; a lateral segmentation of 1 cm (close to one Moliere radius ϱ M 9 mm ) in order to sample with sucient accuracy the lateral shower development (typical scale λ int ). The CREAM calorimeter is a sampling tungsten/scintillating ber device, with an active area of cm 2. It is comprised of a stack of 20 (99.95% pure) tungsten plates, each 3.5 mm thick, and 20 interleaved scintillator layers. In the CREAMI calorimeter, at the corners of each tungsten plate are machined tabs to allow mechanical mounting onto the honeycomb pallet that acts as the base of the support structure. Each plate is approximately one X 0 in thickness, for a total of 20 radiation lengths. Each scintillator layer is made up of fty 1 cm wide, 0.5 mm thick scintillating ber ribbons. The ribbons of the lowest layer are mounted on an aluminum plate with identical shape and dimension to those of the tungsten plate. Alternate layers are oriented orthogonally to each other,

54 50 The CREAM Detector SCD Graphite Target S0/S1 S2 Calorimetr Figure 2.7: The calorimeter module: SCD, hodoscopes, graphite target and calorimeter. providing measurements in XZ and in YZ view, thereby allowing 3dimensional track reconstruction. It is remarkable to notice that the radiation length of tungsten, at 6.76 g/cm 2, is comparable to that of lead (6.37 g/cm 2 ); but, with a higher density of 70%, tungsten layers are 37% thinner per X 0, minimizing the height of the stack and thereby increasing the geometric acceptance. The CREAMII calorimeter was built by the INFN group of Siena/Pisa and completed in July The baseline mechanical design has been modied to allow for a modularization of the stack into completely independent lighttight elements. Each tungsten plate is glued onto a carbon ber composite frame which accomodates 25 lightguides on each side. Superlayers consist of 4 calorimeter layers each. The disassembly of the calorimeter can be done either into superlayers or XY layer pairs. Modularization has been introduced because of the dicult conditions during payload recovery on the ice (learned lesson after the rst ight). Optical splitting design and Readout scheme In the calorimeter, shower particles, generated in the tungsten absorbers, traverse the berribbons producing scintillation light. Each berribbon is aluminized on one end and glued into an approximately adiabatic Acrylic light mixer on the readout end. Aluminization on the nonreadout side of each ribbon improves light collection in each ribbon, and increases the eective attenuation length in the ribbon, thereby reducing the position dependence for showers. For mechanical reasons, alternate ribbons in the same layer are read out on opposite ends. Each light mixer has a bundle of 48 thin clear plastic bers (each with a 256 µm diameter) with a black polyethylene jacket glued into its other end. The clear ber bundle carries the scintillation light signal to 73pixel HPDs identical to those used to read out the hodoscopes. Each bundle is split into three groups each with a dierent number of thin bers, and through a dierent neutral density (ND) lter (optical splitting). The subbundles are glued into a plastic cookie that holds them in position against dierent HPD pixels (see Figure 2.9(c)). This arrangement forms three

55 The CREAM Detector mm W layers 0.5mm scintillating fibers mm fiber ribbons clear fibers bundle Figure 2.8: W/SciFi calorimeter structure and berribbon bundle. readout ranges with a relative signal strength ratio of about 30 from low- to mid-range, and another factor of 30 from mid- to highrange. The readout design is based on the use of HPDs because of their attractive features: operation at a lower gain (a factor 25 to 0) with respect to PMTs, intrinsic linearity of the diode and single photoelectron detection capability. A pair of IDEAS VA32-HDR2/TA32C ASICs read the signals from each HPD, with the lowrange signals in one ASIC, and the mid- and highrange signals in the other. These IDEAS ASICs have a dynamic range of just under 12 bits. With the 3range readout this range is extended above the requirement, even taking into account a S/N ratio of four at the low end. Overlaps between the ranges allow interrange calibration with shower data. The two opposite ends of each layer are read out with an HPD. Each HPD reads out 25 ribbons, with 25 pixels viewing lowenergy range, 25 reading out midenergy range, and 5 more combining highenergy range signals from 5 nearby ribbons each. A single HPD, with two 32 channel ASICs, reads out the 55 signals, as well as three LED (Light Emitting Diode) signals. Six more pixels are read out with no light source, thereby providing a good handle on collective electronic behavior (pedestal drift, gain drift, etc.) for oline calibration. Nine of the 73 pixels are not read out at all. Five HPDs are powered by a single 12 kv HV power supply and read out with the same motherboard. Two motherboards in the same readout box are used to read out the HPDs that read out one calorimeter side. Four such boxes are used to read out the entire calorimeter Trigger System The ux of cosmic rays follows a powerlaw distribution over many orders of magnitude in energy. The CREAM trigger is required to quickly decide which of the large number of particles, traversing the apparatus during its ight, is above the threshold energy and inside the experiment's geometry. The requirements of the trigger are thus to accept all

56 52 The CREAM Detector Shower Particles Tungsten Light Guide Clear Fiber Hybrid Photo Diode Readout Electronics Scintillating Fiber Ribbon Scintillating light Signal from HPD (a) (b) (c) (d) (e) Figure 2.9: (a) Scheme for reading out signals generated in the calorimeter; (b) CREAMII calorimeter; (c) cookie; (d) scintillating bers, light mixer, clearbers bundle; (e) mounted HPD. relevant information about any good particle (i.e. have high eciency) and to keep the background level (i.e. outofgeometry events and events below the threshold energy) low enough to t within the telemetry bandwidth without introducing any bias. The CREAM trigger system is comprised of a master trigger (see Figure 2.) which is set up to allow independent triggering of the TCD and of the calorimeter module. The calorimeter trigger is designed for high energy shower events, while the TCD trigger is thought for events in which a highcharge nucleus traverses the TCD/TRD.

57 The CREAM Detector 53 TCD Module Calorimeter Module Threshold Settings from Command System TCD Data Readout TCD Trigger Fast Calorimeter Trigger Calorimeter Data Readout Master Trigger & Event Counter to SFC Prescaler Settings from Command System External Trigger to SFC Trigger Path Data Path Light Path to SFC Master Trigger Figure 2.: Trigger signal ow in CREAM. The master trigger generates the signal initiating data collection from the subdetectors and building the data into an event record. Calorimeter Trigger The calorimeter trigger requires activity in N consecutive layers (with N set via the command system to be 4, 6, 8 or ) to generate a high energy shower trigger. The calorimeter trigger board accepts the trigger signals generated from each of the calorimeter HPD boxes, and produces a trigger output to the master trigger whenever the selected consecutive layer requirement is met. The calorimeter HPD boxes provide 40 halflayer signals, which are combined to form 20 layer signals. A layer is considered active if at least one ribbon (out of 50), in the low energy range, registers a signal exceeding a commandable threshold. The rate for this trigger in ight ranged from 0.2 Hz to 3 Hz, depending on altitude and threshold setting. As a result of a malfunction in the readout system, the input from one or more layers can be disabled, those layers can be bypassed (equivalent to alwayson condition) to allow continued triggering. The calorimeter trigger accepts commands from the command system to modify a layer bypass mask and/or the number of consecutive layers required for the high energy shower trigger. TCD Trigger The TCD trigger uses a redundant set of coincidences between several PMT signals. The TCD provides a valid trigger signal when a charged particle has traversed both scintillator planes. Combining these signals, the TCD trigger generates one of several trigger ags as appropriate. A highz trigger is sent when a particle of Z Z 0 (e.g. Z 0 = 3)has traversed the detector, while a lowz trigger is sent when a particle of Z 1 passed the TCD and a conrming calorimeter trigger signal arrives within 0 µs. A single side of the TCD is triggered by a logical OR of the four PMTs reading out signals from that side. The lowz trigger ag is generated when a threshold is crossed on the anode signal of at least one PMT on each of at least three sides of the TCD; the same

58 54 The CREAM Detector condition is required to be also met for at least one side of S3. A highz trigger signal is generated when a threshold is exceeded on the dynode 11 signal of at least one PMT on each of at least three sides of the TCD; the CD is required to have recorded a signal above a threshold corresponding to an incident particle with β 0.7. This last condition reduces the background from the plentiful low energy highcharge nuclei expected above Antarctica, where the geomagnetic cuto is much lower than in equatorial latitudes. 2.3 Balloncraft The term ballooncraft is used to designate all hardware below the attachment point to the mobile launch vehicle. This hardware, with the exception of a rotator, forms an integrated assembly containing the instrument and support systems mounted on the primary support structure. As shown in Figure 2.11, the ballooncraft consists of the ballooncraft support system and the instrument package. The ballooncraft support system has the capability to provide global coverage of telemetry downlinks and command uplinks, power, and data processing for the science instrument packages. In fact,originally designed to y as part of NASA's Ultra Long Duration Balloon (ULDB) program, with ight durations of 600 days, and with trajectories expected to cover much of the southern hemisphere (30 C South latitude down to the South Pole), the CREAM payload has several properties not shared with any other payload own before in the Long Duration Balloon (LDB) program. These include the use of a highgain antenna to provide near realtime data downlink, and command uplink, through the Tracking and Data Relay Satellite System (TDRSS), as well as continuous control of the payload. The ballooncraft support system consists of the Command and Data Module (CDM), instrument support structure (ISS), solar array panel support system, antenna support systems, ballast hopper and crush pads. It also monitors the health and well being of the balloon vehicle and the ballooncraft. The CDM is the core of the ballooncraft support system and contains the control and monitoring of the ballooncraft power system and communications through a ight computer subsystem. It also controls and monitors the health and safety of the overall balloon system, archives data, and provides balloon positional data. The CDM is located in an independent thermally controlled support structure. Mechanical system The ISS provides the primary mechanical support for both the instrument and the ballooncraft support system. The mechanical structure and associated hardware are designed and built to meet the strength requirements mandated by the National Scientic Balloon Facility (NSBF). The ballooncraft is designed to survive a nominal ground impact. It is recognized, however, that some structurally mounted hardware such as the solar array and TDRSS antenna, may not survive a nominal impact. The parachute is able to automatically separate from the ballooncraft upon ground impact. The ballooncraft is designed such that all instrument electronics boxes, calorimeter, carbon blocks, TRDs, hard drives and TDRSS transponder can be removed with minimal tools during recovery.

59 The CREAM Detector 55 CREAM Instrument CDM DATA system Ballast Hopper Crush Pads Figure 2.11: CREAM ballooncraft conguration. Solar Array Wing The CREAM block diagram in Figure 2.12(b) shows the interface between the science instrument and the ballooncraft support system. Downlinked data is received by White Sands and forwarded to the Operations Control Center (OCC) in Palestine, TX. From there the data is sent to the Engineering Support Center (ESC) at NASA's Wallops Flight Facility (WFF) on Wallops Island, VA, and on to the Science Operations Center (SOC) at the University of Maryland (UMD), as shown in Figure 2.12(a). The instrument can communicate with the Remote Operations Control Center (ROCC) at Williams Fields, outside of McMurdo Station, Antarctica, through a Line of Sight (LOS) transmitter from before launch until the payload has gone over the horizon about 12 hours post launch. During the ight, monitoring and archiving data at the SOC are performed by the Linux based CREAM Data Acquisition software (CDAQ). Support functions for the CREAM payload (e.g. power, data archiving, GPS position determination, command link, telemetry, etc.) are handled by the WFF CDM. When the payload enters a TDRSS Zone of Exclusion (ZOE), the data downlink is restricted to a reduced housekeeping packet every 15 minutes via the backup IRIDIUM link. Once the payload has completed its ZOE traverse, the CDM plays back the backlog of archived data, in parallel to the thencurrent data telemetry. With a total TDRSS downlink bandwidth of 85 kilobits per second (kbps), 35 kbps is used for the playback until the backlog is cleared, with a near realtime data rate of about 50 kbps. During the ight, TDRSS command windows as long as 23 hours are made available. All high energy events, in which a signicant shower is recorded by the calorimeter, are transmitted to the SOC via TDRSS, along with % of events in which a particle with Z 3 traversed the TCD. Weight and Power Two major considerations in balloon experiments are the limited weight the balloon can carry to oat within a certain range of altitude and the limited power that can be provided by either batteries or solar arrays. The CREAM instrument is designed to optimize both

60 56 The CREAM Detector Iridium ARGOS GPS TDRSS Iridium ARGOS e mail WSCT SOC Iridium Gateway ESC OCC ROCC TDRSS LOS Backup (a) Trigger Electronics Charge Detector TRD Calorimeter Module Optical Readout Front End Electronics Optical Readout Data Data Data CREAM Science Flight Computer Commands & Data Ethernet Housekeeping Sensors Command Processor LV Distribution Photomultiplier HPDs High Voltage Power Supplies Low Voltage Power Supplies 28 VDC Ballooncraft Interface (b) Figure 2.12: (a) CREAM ight operation scheme. (b) CREAM block diagram.

61 The CREAM Detector 57 the weight and power while collecting high quality data at a rate as high as possible. The 40 million cubic foot light longduration balloon used by the NASA allows a suspended weight of up to 2727 kg. After considering the ight train (cables, parachute), support systems (power, telemetry, rotator, crush pads), an instrument weight of about 1143 kg (see Table 2.1) allows 500 kg (750 kg) of ballast. The CREAM power system includes a solar array for electrical power generation and batteries for energy storage. The power system is designed to provide sucient power to support a nominal 0day mission with 0% duty cycle. The power system is sized for nominal ight operations within the range of 90 south latitude to 30 south latitude. The nominal operating output voltage of the power system is 28±4 Vdc. With a panel solar array, the CREAM power system has been able to generate about 1400 W during its rst ight. The support systems used an average of about 400 W, while the instrument about 380 W (345 W). For detail, see Table 2.1. Subsytems Weight [kg] Power [W] CREAMI CREAMII CREAMI CREAMII Upper detectors Targets & interleaved detectors Calorimeter Common electronics Total CREAM instr Ballooncraft sup. sys Table 2.1: Payload weight and power summary for CREAMI and CREAMII. Thermal system The thermal system is designed to maintain a high level of reliability in the electronics, and to maintain the integrity of the pressurized gas in the TRD makeup reservoir. The required temperatures are: (i) C to 50 C for the electronics (ii) 20 C to 50 C for the calorimeter and charge detector (iii) 20 C to 40 C for the TRD (iv) 0 C to 40 C for the batteries. The thermal design utilizes an isothermalized basepanel, with conventional heatpipes connected to radiator panels (also doubling as shearpanels). To provide heat rejection at high temperatures while retaining heat at low temperatures, CREAM utilizes a heat switch in the form of motorized shades. Above C the shades are fully open, while below 0 C they are fully closed. The shade material is specially designed and tested to provide sucient thermal insulation to keep the payload at acceptable temperatures even

62 58 The CREAM Detector with extremely cold ambient temperatures. The upper modules are radiatively coupled to the basepanel with an aluminized thermal blanket covering the top of the payload. The thermal switch technique is a simple, inexpensive and elegant solution that allows the payload to remain warm enough in cold environments and cool enough in warm environments, i.e., to remain fully operational, without resorting to hundreds of watts of heater power.

63 Chapter 3 CREAMII ight and performance Redundancy in the direct measurement of the electric charge of cosmicray nuclei is provided in the CREAMII instrument by the SCD silicon arrays, the TCD scintillator paddles and the Cherenkov counter, while the tungstenscintillator sampling calorimeter is used to measure their energy. The preliminary analysis presented in this work is made using only the information provided by 2 subsystems, the SCD and the calorimeter. In this chapter a summary of their ight performance is reported. 3.1 The second ight of CREAM The CREAM payload was successfully own twice on longduration balloons from Mc- Murdo, Antarctica in 2004/05 and 2005/06. The second launch occurred on December 16 th 2005, exactly 1 year after the rst launch. This ight circumnavigated the Pole twice before it was terminated on January 13 th As shown by the trajectory of the second ight (Figure 3.1), the balloon initially drifted east and then looped southward and westward. It was visible when it came back to McMurdo after the rst circumnavigation. It spiraled out northward and spent a few days over the water during its second circumnavigation. The ight was terminated when it came back to North of McMurdo. The payload landed 290 nautical miles northwest of McMurdo station with minimal damage. It was much closer than the landing site of the rst ight, so its recovery was much easier. The position of the payload was monitored during the whole ight using the Global Positional System (GPS). The GPS dispositive gave information on the altitude as well as of the latitude and longitudinal position. The GPS data was used both to track the payload and to predict possible landing places. The balloon oat altitude was between 35 and 40 km throughout most of the ight, as shown in Figure 3.2(a). The corresponding average atmospheric overburden was only 3.9 g/cm 2. The diurnal altitude variation due to the Sun angle change was very small, < 1 km, near the Pole, i.e. at high latitude, although it increased as the balloon spiraled outward to lower latitudes. The temperature of the various instrument boxes stayed within the required operational range with daily variation of a few C, consistent with the sun angle. All highenergy data were transmitted via TDRSS during the ight, while the lower energy data were recorded on board. A total of 57 GB of data including 27 6 science events were collected from the second ight. 59

64 60 CREAMII ight and performance Figure 3.1: Balloon trajectory of the second ight. The red curve represents the rst circumnavigation, the green one the second. 3.2 Trigger logic A detailed description of the trigger system and its implementation during the second ight is beyond the scope of this paper (see Figure 3.3). For a discussion of the preliminary analysis reported herein, it is sucient to consider four main trigger ags. The CAL trigger discriminates against the abundant lowenergy cosmicray particles, by requiring the detection of a shower in the calorimeter. The topological cuts dening the minimal requirements for a shower are the presence of at least 6 consecutive calorimeter layers, each with at least one hit above a given threshold. During most of the ight, the latter was set around a value equivalent to 60 MeV energy deposit in one ribbon. The CAL & ZLO and CAL & ZHI trigger ags require, as an additional condition to the CAL trigger, the presence of a scintillation signal detected by at least 3 PMTs of the TCD, compatible with the one expected from a lowz or highz nucleus, respectively. The threshold to discriminate a lowz and a highz nucleus was set respectively at 0.5 (i.e. below the proton peak) and 1.5 equivalent particle charge (i.e. above the proton, but below helium peak). The ZCLB trigger ag occurs when the calorimeter is not triggerred and a scintillation signal, consistent with a highz nucleus (Z 2), is detected in at least 3 TCD PMTs in coincidence with a signal from the S3 detector. In Table 3.1 the conversion of the trigger logic signals into the above mentioned trigger ags, useful for the present analysis, is reported.

65 CREAMII ight and performance 61 (a) (b) (c) Figure 3.2: (a) Balloon altitude. (b) High voltage versus time for all calorimeter's modules. (c) Operative temperature for the electronics crates on the sunside (black curve) and antisun side (red curve).

66 62 CREAMII ight and performance ZLO AND ZLO*S3HI Prescale Low_E Trigger ZHI ZHI Prescale S3LO ZHI*S3LO ZCLB S3HI S3HI CAL*S3HI Prescale OR MT CAL CAL High_E Trigger EXT CALIB Beamspill AND Beamtest External Beamtest Calibration Pedestal Figure 3.3: CREAMII master trigger logic. Trigger Type Trigger Logic Trigger Number CAL EHI 4288 EHI & ZLO 4292 CAL & ZLO EHI & ELO 6340 EHI & ZHI 4296 EHI & ZHI & ZLO 4300 CAL & ZHI EHI & ZCLB & ZLO 5324 EHI & ELO & ZHI 6348 EHI & ELO & ZCLB 7372 ZCLB ZCLB 1152 Table 3.1: CREAMII trigger ags. During the ight, diurnal variations of the CAL trigger rate were recorded with peak values observed in correspondence of the lowest balloon altitudes. The more restrictive CAL & ZLO and CAL & ZHI triggers did not show signicant rate variations during the daynight cycle. Their average rates were Hz and Hz respectively.

67 CREAMII ight and performance Livetime The detector livetime is dened as the percent of time during which the detector is ready to capture an event. In CREAMII ight, a pair of timers are used to measure live and total time. These counters are incorporated into the housekeeping (HK) system and a trigger live signal is used from the Master Trigger Board to measure the livetime. The counters are each 48bit, with a khz clock to provide a resolution of about 3.25 µs. Since only an average bandwidth of 0.5 kbps was assigned for telemetry, each housekeeping event was recorded and downloaded every 6 second. Three temporal information are recorded for every housekeeping event: the totaltime T t, the livetime T l and the event time T evt which is expressed in GMT. For each HK event, T l vs. T evt values can be plotted: a typical behaviour of the livetime information over a day of ight is shown in Figure 3.4(a). It can be noticed that livetime increases for increasing T evt values, as it is expected, but two general anomalies can be recognized: sometimes the livetime values dropped during the ight to a lower value; after this drop, the usual increasing behaviour as a function of ight time can be seen; there are some housekeeping events whose recorded T l value is higher or lower than the expected one from the general shape of the prole plot of T l vs. ight time. A specic algorithm has been developed for correcting for these anomalies and evaluating separately the totaltime and the livetime of the experiment. Then the consistency between T t and T l values is checked on an eventbyevent basis. In fact, for each housekeeping event, the following relationship is expected to be satised: which means δ evt δ t δ l (3.1) = (δ t δ l ) 0 (3.2) where δ evt is the time increase between two consecutive HK events; δ t and δ l are the correspondent totaltime and livetime increase respectively. The distribution of livetime increase δ l and are shown in Figure 3.4(b) and 3.4(c) respectively. The latter shows an anomalous behaviour because of the presence of negative values: in 1% of events the totaltime increase is less than the respective livetime increase. These events are removed from the present analysis. The reconstructed livetime information is shown in Figure 3.4(d) as a fuction of ight time. The preliminary analysis of the second ight data reported in this work is performed on a sample of good runs taken in the period from December 19 th to January 12 th, excluding the following temporal intervals: 6 hours on December 22 nd when a calorimeter motherboard suered a failure; 24 hours from January 3 rd to January 4 th because of the presence of an extra power current in the science ight computer;

68 64 CREAM II ight and performance [min] :00 20/12/05 05:00 20/12/05 08:00 20/12/05 11:00 20/12/05 14:00 20/12/05 17:00 20/12/05 20:00 20/12/05 23:00 20/12/05 (a) Counts Counts δl [sec] Livetime [min] (b) δl and [sec] (c) /12/05 22/12/05 26/12/05 02/01/06 29/12/05 05/01/06 09/01/06 13/01/06 (d) Figure 3.4: (a) Live time information over a single day of ight. Distributions of values for all the ight period are shown in (b) and (c) respectively. (d) Reconstructed live time as a function of ight time.

69 CREAMII ight and performance hours on January 6 th because of an anomalous increase in noise in ZHI triggers. The totaltime for the selected period amounts to minutes, while the eective livetime is estimated to be minutes, providing an average of 75%. 3.4 Calorimeter performance High voltage and thermal behaviour The hybrid photodetectors of the calorimeter require an operational voltage of several kv. The system has to operate at low atmospheric pressure (a few mbar) as the payload is not pressurized. During the whole ight, the calorimeter high voltage system performed successfully with an average voltage set at 7 kv. An exception was the HV trip of one module on December 23 rd (Figure 3.2(b)). The high voltage was promptly restored at a lower voltage of 6 kv and nally set to the original value. A set of temperature sensors monitored the calorimeter's thermal behaviour, which was found in good agreement with the expectations of the ight thermal model. The operative temperature range was between 20 and 35 C for the electronics crates located on the sunside, and 15 to 30 C for those on the antisun side (Figure 3.2(c)). A daily temperature uctuation of about C was observed, depending on the inclination of the sun light during each 24hour cycle Frontend electronics behaviour The behaviour of the photodetectors and frontend electronics of the calorimeter was monitored during the whole ight, by means of special calibration runs repeated every two hours: the HPD response stability was checked by illuminating a few pixels with LED light; the electronics channels variations were monitored by injecting a test charge. average pedestals rms [ADC] day of flight (a) (b) Figure 3.5: (a) Pedestal values of the channels of one HPD vs. day of ight. (b) Calorimeter average pedestal rms noise during the ight.