Workshop on AstroParticle Physics, WAPP 2009Bose Institute, Darjeeling, December 2009

Extensive Air Showers and Astroparticle Physics Observations and Interpretation through Simulations

Suresh TonwarDepartment of Physics, University of Maryland,

College Park, MD 20742, USA

Extensive Air Showers and Astroparticle Physics Observations and Interpretation through Simulations

Plan

1. Historical Remarks2. Introduction to Primary Cosmic Rays3. Development of Extensive air showers4. Observable in Extensive Air Showers5. Measurement of Observables6. Interpretation of Observables7. Results in AstroParticle Physics

PRIMARY COSMIC RAYS

• Low energy cosmic rays – 10 MeV – 10 GeV

Protons, He, C, N, O, …., Ne, .., Si, .. Ca, … Ni, Fe, up to the highest Z, including transuranic elements Electrons and positrons Antiprotons and antinuclei (??) Photons Neutrinos

Particle identification and energy spectrum, directional distribution

Solar cosmic rays, Solar flares, Magnetic disconnections and Heliospheric Modulations Galactic cosmic rays, Acceleration by supernova shocks, Neutron stars (pulsars) Stellar Nucleosynthesis Galactic propagation, mean life time, energy budget

Detailed observations (elemental and isotopic) with detectors flown aboard balloon and satellite borne platforms Magnetic spectrometers, scintillation detectors, Cherenkov radiation detectors, Transition radiation detectors, Silicon strip and pixel detectors

PRIMARY COSMIC RAYS

• High energy cosmic rays - 10 GeV – 1000 GeV

Protons, He, C, N, O, …., Ne, .., Si, .. Ca, … Ni, Fe Electrons and positrons Antiprotons, Antinuclei (??), Dark Matter (??) Photons Neutrinos

Particle identification and energy spectrum, directional distribution

Galactic sources, galactic confinement by galactic magnetic field (~ microgauss) Acceleration by supernova shocks, neutron stars

Direct observations (elemental) with detectors flown aboard balloon and satellite borne platforms

Electromagnetic and hadronic calorimeters, scintillation detectors, Cherenkov radiation detectors, Transition radiation detectors, Silicon strip and pixel detectors

Observations with detectors placed at various atmospheric levels, mountain altitudes, sea level and underground (muons and neutrinos)

PRIMARY COSMIC RAYS

• High energy cosmic rays - 10 GeV – 1000 GeV

Protons, He, C, N, O, …., Ne, .., Si, .. Ca, … Ni, Fe JACEE, RUNJOB, …

Electrons and positrons Antiprotons, Antinuclei (??), Dark Matter (??)

PAMELA, AMS

PhotonsMAGIC, HESS, VERITAS, HAWC

NeutrinosAMANDA, ICECUBE, NESTOR,

PRIMARY COSMIC RAYS

• Very High energy cosmic rays - 103 GeV – 105 GeV

Protons, He, C, N, O, …., Ne, .., Si, .. Ca, … Ni, Fe Electrons, positrons and antiprotons Photons Neutrinos

Particle identification and energy spectrum, directional distribution

Galactic sources, galactic confinement by galactic magnetic field (~ microgauss) Acceleration by supernova shocks, neutron stars

Direct observations (elemental) with detectors flown aboard balloon and satellite borne platforms Observations at Antartica

Electromagnetic and hadronic calorimeters, scintillation detectors, Cherenkov radiation detectors, Transition radiation detectors, Silicon strip and pixel detectors

Detectors placed at mountain altitudes, sea level and underground – large area scintillation detectors, resistive plate chambers and Cherenkov detectors – Extensive Air Showers

PRIMARY COSMIC RAYS

• Very High energy cosmic rays - 103 GeV – 105 GeV

Protons, He, C, N, O, …., Ne, .., Si, .. Ca, … Ni, FeTibet AS/Gamma, ARGO-YBJ, GRAPES,

Electrons, positrons and antiprotonsAMS

PhotonsPachmarhi, Mount Abu, MILAGRO

NeutrinosAMANDA, ICECUBE, NESTOR,

PRIMARY COSMIC RAYS

• Ultra High energy cosmic rays - 105 GeV – 109 GeV

Protons, He, C, N, O, …., Ne, .., Si, .. Ca, … Ni, Fe Photons

Neutrinos

Particle identification and energy spectrum, directional distribution

Galactic sources, galactic confinement by galactic magnetic field (~ microgauss) Acceleration by supernova shocks, stellar winds

Direct detection impractical particle flux at 106 GeV ~ 1 m-2 day-1

Extensive air showers Detectors placed at mountain altitudes, sea level and underground – arrays of scintillation detectors, resistive plate chambers and Cherenkov detectors spread over large area to collect sufficient number of particles over months/years

PRIMARY COSMIC RAYS

• Ultra High energy cosmic rays - 105 GeV – 109 GeV

Protons, He, C, N, O, …., Ne, .., Si, .. Ca, … Ni, FeKASCADE, CASA-MIA, Tibet AS/Gamma, GRAPES,

PhotonsTibet AS/Gamma, GRAPES,

NeutrinosICECUBE, NESTOR,

EXTENSIVE AIR SHOWER• A primary cosmic ray particle/photon incident on the Earth from the outer space,

interacts with an air nucleus near the top of the atmosphere, with a mean free path which is different for different particles and is energy-dependent.

• Interaction of a strongly interacting primary (such as a proton or a heavier nucleus) produces strongly interacting secondaries, such as pions, strange particles (K’s), charmed particles and other particles containing heavier quarks.

• Interaction of a high energy photon creates mostly an electron-positron pair while a high energy electron mostly radiates a photon when it interacts with an air nucleus.

• Interaction of a high energy neutrino is very rare in the atmosphere due to its very small interaction cross-section. The interaction, if it occurs produces an electron, muon or a tau meson depending on the flavour of the neutrino. Cosmic ray neutrinos are detected mainly through their interaction with rocks in experiments located deep underground or underwater.

EXTENSIVE AIR SHOWER• In hadronic interactions, secondary particles are produced sharing energy of the

primary particle which mostly continues with reduced energy (inelasticity)

• Strongly interacting (hadrons) secondaries and the surviving primary interact again lower down in the atmosphere, producing more secondaries with energies reduced further. The multiplication process continues till the energy of a particle becomes less than the minimum required to create a secondary particle. All charged particles continue to lose energy continuously by excitation and ionization at the rate of about 2 MeV per gm/cm2 as they travel through the atmosphere.

• Some unstable secondary particles, such as pions, kaons, etc, decay in flight, depending on the relative probabilities for interaction and decay and their Lorentz factor. For example, a 14 GeV charged pion decays to a muon and 2 neutrinos at a mean distance of ~800 m from its production point if it does not interact before the decay point.

• Neutral pions produced in these interactions decay almost instantaneously to two photons soon after their production and start their own electron-photon cascade as the shower of particles travels down the atmosphere.

EXTENSIVE AIR SHOWER

• For electron or photon primaries, the pair production and the radiation processes alternate producing an increasing number of electrons and photons as they travel down the atmosphere.

• In a small number of electron or photon interactions with air nuclei, photo-production of pions and other hadrons also takes place. These pions may produce muons through their decay.

EXTENSIVE AIR SHOWER

• While the longitudinal development of the shower continues along the axis defined by the direction of the primary cosmic ray particle, the shower also spreads laterally due to the transverse momenta acquired by secondary hadrons at production. Interestingly the distribution of the transverse momentum for various hadronic secondaries is very similar, ~pt exp (-pt/p0) and the average value is also not very different for pions, kaons and baryons.

• Due to this transverse momentum, the hadronic secondaries spread around the shower axis to several tens of meters. The electrons and photons produced by these hadrons spread out to much larger distances laterally due to scattering with air atoms.

EXTENSIVE AIR SHOWER

• Finally, as the shower particles reach the observational level, either at mountain altitude or sea level, there are hadrons, muons, electrons, photons and neutrinos of various energies.

• Measurements of various properties of these particles, observed with arrays of different type of particle detectors, are then used to determine the basic properties of the primary cosmic ray particle, such as its energy, mass and arrival direction (zenith angle and azimuth angle)

E.A.S. OBSERVABLES ELECTRONS & PHOTONS

• Due to the transverse momentum acquired by neutral pions and other hadrons at production and due to scattering of the electrons by air atoms, the electron-photon component is spread over a large distance around the direction of incidence of the primary particle at the observational level. The distribution around the expected position (x0, y0) of the shower core is called the LONGITUDINAL DISTRIBUTION of the electron component. It is approximately exponential function of the core distance (re), exp (-re /r0), where r0 is characteristic of the observational level in the atmosphere. For example, r0 ~ 105 m for the Darjeeling level and 90 m for the Siliguri level..

• The total number of electrons and photons, called Shower Size Ne, in an EAS is nearly proportional to the energy of the primary particle transferred to neutral pions and other particles which decay to neutral pions and/or photons. On the average, Ne ~ E0

1.1 i, in the lower atmosphere, where E0 is the energy of the primary particle. The proportionality constant depends on the type of primary particle and the observational level.

E.A.S. OBSERVABLES ELECTRONS & PHOTONS

• A measurement of the shower size gives a reasonably good estimate of the average energy of the primary particle, provided we know the nature of the particle, e.g., proton, He nuclei, ..,or photon. However, large fluctuations in the development of individual showers results in a large spread in the mean energy determined from the observed shower size for a given primary energy.

• Observationally, the shower size can be measured by laying a carpet of particle detectors at the observational level. However, the carpet area has to be necessarily in thousands or tens of thousands of square metres, which is prohibitively expensive. The required carpet area increases with increasing primary energy. Presently, only 2 experiments have tried this approach for measuring the primary energy, ARGO-YBJ at Tibet and MILAGRO at Los Alamos in USA.

E.A.S. OBSERVABLES ELECTRONS & PHOTONS

• More commonly, the shower size is estimated from a sampling of particle densities at various distances from the shower axis, using a moderately large number of particle detectors spread over an area ~ 105 – 109 m2, depending on the primary energy of interest. The accuracy in the measurement of shower size improves with an increase in the number of detectors available for measuring the lateral distribution of electrons and photons.

• Typically, a small array like the 50-detector array installed by the TIFR group above the L3 detector at CERN, permits the measurement of the mean energy with an accuracy of about a factor of 2, covering an area of only ~ 3000 m2. An array like the 350-detector array at Ooty, the GRAPES-3 array, permits an accuracy of about 30% for primary energies around 106 GeV, covering an area of ~ 30,000 m2. The very large array at the Pierre Auger Observatory in Argentina, covering an area of 3000 km2 with 3000 detectors permits the measurement of energy to ~10% accuracy for showers initiated by protons of 1010 GeV.

E.A.S. OBSERVABLES ELECTRONS & PHOTONS

• In addition to the estimate of shower size required for estimating the primary energy, a measurement of the arrival direction of the primary particle is also necessary for interpreting the observations.

• Determination of arrival direction (zenith angle theta and azimuth angle phi) is made using the relative time differences between the arrival of particles at various detectors. Since most of shower particles are produced several kilometres above the observational level, it is convenient to assume the shower particle disk (about 2-3 m thick) to be a plane.

• It is common to measure the space angle of the shower direction to an accuracy of 0.60-1.00 depending on the number of detectors available for reconstruction of the shower direction. Due to scattering of electrons resulting in fluctuations in their arrival time over different detectors, shower arrays with compact configurations, like the GRAPES-3 array, provide better accuracy. This is very helpful when searching for discrete astrophysical sources in the galaxy for cosmic ray photons.

E.A.S. OBSERVABLES ELECTRONS & PHOTONS

• Typically, a small array like the 50-detector array installed by the TIFR group above the L3 detector at CERN, permits the measurement of the mean energy with an accuracy of about a factor of 2, covering an area of only ~ 3000 m2. An array like the 350-detector array at Ooty, the GRAPES-3 array, permits an accuracy of about 30% for primary energies around 106 GeV, covering an area of ~ 30,000 m2. The very large array at the Pierre Auger Observatory in Argentina, covering an area of 3000 km2 with 3000 detectors permits the measurement of energy to ~10% accuracy for showers initiated by protons of 1010 GeV.

40-detector array at P2 (LEP/LHC, CERN)

40-detector array at P2 (LEP/LHC, CERN)

20-detector array at P4 (LEP/LHC, CERN)

20-detector array at P4 (LEP/LHC, CERN)

GRAPES-3 Array at Ooty

• GRAPES-3

E.A.S. OBSERVABLES MUONS

• The muons are spread over much larger distances laterally around the shower axis as they travel more or less directly from the point of production to the observational level, without much scattering. Lower energy muons lose energy mainly by ionization while higher energy muons also lose energy occassionally by radiation (bremsstrahlung). A small number of lower energy muons are also lost during their passage in the atmosphere due to decay.

• Muons are identified after filtering out particles which have strong or electromagnetic interactions like hadrons and electrons by placing thick absorbers above the particle detectors. Typically, we need to put 600 gm/cm2 of absorber (~2.5 m thick concrete block) above the particle detector to identify muons of energy larger than 1 GeV, as in the 640 m2 muon detector of the GRAPES-3 array at Ooty.

E.A.S. OBSERVABLES MUONS

• The L3 muon spectrometer observed muons of energy larger than 15 GeV as the spectrometer was located 30 m underground, near Geneva.

• A TIFR group operated a 20 m2 detectors 220 m underground in the Kolar Gold Mines in 1980’s to study muons of energy larger than 220 GeV in air showers detectors with a 127-detector array above the surface.

• The Proton Decay group at TIFR operated a detector at 2 km underground in KGF mines to look for TeV muons in 1980’s.

• The MACRO group in Gran Sasso Laboratory in Italy studied muons of energy larger than 1.6 TeV using the mountain above as the filter in 1990’s.

E.A.S. OBSERVABLES MUONS

• The total number of muons, Nmu, of energy greater than some low energy threshold, say 1 GeV, is a good measure of the energy of the primary cosmic ray particle, if its nature is known. Alternatively, the muon size, Nmu, can be used to determine the nature of the primary particle since Nmu depends sensitively on the atomic number of the primary nucleus.

• Measuring Nmu is a difficult task as the muons are spread over much larger area than the hadrons or the electrons. Therefore, most of the experiments have attempted to sample muon densities at a few locations in the shower array.

• The GRAPES-3 experiment samples the density of muons of energy larger than 1 GeV at a distance of ~ 60 m from the centre of the shower array, using the 640 m2 area tracking muon detector.

• The KASCADE group operated a large number of muon detectors to sample the muon density under each electron detector to estimate the muon size for muons of energy larger than 250 MeV. In addition, the collaboration also operated a 200 m2 detector at the centre of the array to observe muons of energy larger than 1 GeV.

E.A.S. OBSERVABLES HADRONS

• Most of the high energy (Eh > 10 GeV) hadrons (protons, antiprotons, neutrons, antinueutrons, charged pions, kaons, etc.) are located very close to the shower axis, mostly within 20 m. It is rather difficult to measure the energy spectrum or the number of high energy hadrons in a shower as the cascades produced by various hadrons overlap in the calorimeters used to detect and measure the hadrons. The large multiplate cloud chamber, operated at Ooty in 1970’s, was the only instrument which permitted the measurement on individual hadrons at mountain altitude and provided a unique measurement on the ‘Neutral/Charged’ ratio for high energy hadrons. This ratio is sensitively related to the baryon production cross-section in high energy hadron-air collisions at very high energies.

• The number of high energy hadrons in a shower decreases rapidly with increasing atmospheric depths. Showers of energy ~ 106 GeV have very few hadrons of energy larger than 100 GeV at the sea level. Practical difficulties with flux detectable with small area calorimeters has prevented detailed studies on high energy hadrons in air showers.

E.A.S. OBSERVABLES Cherenkov Photons

• All charged particles of the shower travelling with velocity larger than the velocity of light in the air emit Cherenkov photons which can be detected with arrays of optical detectors placed at mountain altitudes or sea level. The number of Cherenkov photons in a shower is a very good measure of the energy of the primary particle. However, these photons are spread over large distances, typically over 150 m around the shower axis. A very large number of optical detectors are required to sample the photon densities over the large area.

• Though the number of Cherenkov photons is a good energy estimator for the shower, the technique has not been used widely due to the serious inefficiency associated with the observations. These studies can be carried out only during clear moon-less, fog-less nights which are available only during ~ 8% of the time, even at the best observational sites such as at Hanle in Ladakh or in the Namibian desert (HESS telescope).

Data Collection : Jun 2004 – Dec 2006• 2004 (174 days); 2005 (186 days);2006 (350 days)

• Total 710 days

• Data analyzed for each day separately to check for various distributions, like event rate, inter-event time separation, ADC and TDC distributions, etc.• ADC data converted to MIPS and ‘Particle Sum’ calculated – approx estimator for primary energy• TDC data used to calculate shower arrival angle (zenith and azimuth)

Shower rate variation over 24 hours

Particle sum spectrum at P2

TDC (ns) distributions at P2 w.r.t D01

P2-P4 Shower Arrival Time Distribution, Jan 1- Mar 1, 2006Max. time difference expected is ~ 8 km cos(45)/c ~ 20 ns

P4-P2 Shower Arrival Time Distribution, Jan 1- Mar 1, 2006

P2-P4 Shower Arrival Time Distribution, Jan 1- Mar 1, 2006

P2-P4 Shower Arrival Time Distribution, Jan 1- Mar 1, 2006

P2-P4 Showers Space Angle Distribution, Jan 1- Mar 1, 2006

P2-P4 Showers Space Angle Distribution, Jan 1- Mar 1, 2006

Present status of data analysis – Jun 30, 2009

• 220 days data analyzed

• 554 showers (with info on arrival angle) observed with time separation < 30 us• 525 expected from chance• Good agreement between observations and expectations

• No shower pair observed as yet with space angle less than 8 degrees and time separation less than 20 us

Conclusions and Outlook

• No time and arrival angle coincident (dtime < 20 us, dtheta < 5 degrees)

shower pair observed from an analysis of data for 220 days, out of available data for 710 days

• Data analysis expected to be completed by the end of 2009

• If no signal is observed after the completion of the analysis of the full data set, an upper limit on the flux of ‘strangelets’ can be calculated

THANKS FOR YOUR KIND ATTENTION