to

1

• Introduction
• to
• Cosmic Rays
• Dr. Johana Chirinos
• Michigan Tech University

2
Cosmic Rays

• Discovery by Hess in 1912.
• Cosmic rays
• high energy particles incident on Earth
• from outer space,
• plus secondary particles,
• which they generate as they traverse atmosphere.
• Their study
• pioneering role in study of elementary particles
• and their interactions.
• -Discovery in cosmic rays of antimatter
• positron in 1932
• -Discovery of ps and µs and strange
• particles in 1940s
• kick-started building of large particle
  accelerators and
• development of detection equipment essential for
  elementary particle physics.

3
Cosmic Rays

• In 1990s study of interactions of solar and
  atmospheric ?, on distance scales far larger
• than anything at accelerators or reactors,
  revealed first cracks in Standard Model
• evidence for ? flavour mixing and for finite ?
  masses.
• This led to a revival, in the new millenium,
• of lepton physics in fixed-target experiments at
  accelerators,
• development of new proposals
• building of µ storage rings to serve as sources
  of high energy ?e and ?µ.
• At highest energies, studies of ?-rays in TeV
  range and above indicated point sources in
• the skies where it seems the most violent events
  in the universe have taken place.
• ?ntensive studies of both ?-rays and ultra-high
  energy protons and heavier nuclei will
• shed new light on mechanisms for particle
  acceleration, as well as new fundamental
• processes at energies far in excess of what could
  ever be achieved on Earth.

4
The spectrum and composition of cosmic rays

• Charged primary particles
• - protons (86)
• - a-particles(ll)
• - nuclei of heavier elements up to uranium (1),
  
• - electrons (2).
• ?lso very small proportions of positrons and
  antiprotons, of secondary origin,
• generated by interactions of primary particles
  with interstellar gas.
• Neutral particles ?-rays, ? and anti-?.
• Some can be identified as coming from point
  sources in the sky like
• -? from Sun and supernovae,
• -? rays from sources such as Crab Nebula and
  active galactic nuclei(AGN).
• Auger Observatory charged primary correlated
  with AGN

5
The spectrum and composition of cosmic rays

• Direct
• Primary particles identified with nuclear
  emulsion detectors in high altitude balloons.
• Detectors flown in satellites
• -scintillation counters to measure primary
  nuclear charge.
• -gas-filled Cerenkov counters to measure particle
  velocity and energy.
• Indirect
• For GeV-TeV energy region, calorimetric method
  has been employed
• Measuring ionization energy in electromagnetic
  showers that develop as a result of
• nuclear cascade which the primary generates as it
  traverses the thicknesses of absorber.

6
Energy spectrum of cosmic ray protons

• Above energy of few GeV up to
• so-called knee at 1014 eV,
• spectrum follows a simple power law
• N(E) dE const. E-2.7 dE
• Above knee, spectrum becomes steeper
• with an index of about -3.0,
• before apparently flattening off again
• above 1018 eV at ankle.
• 1020 eV
• 1 baseball(140g) at 60mph
• Acelerators lt 1015 eV

7
The spectrum and composition of cosmic rays

• Primary spectrum multiplied by E2.7, showing the
  knee.
• Isotropy
• At energiesgt30 GeV, where effects of magnetic
  fields of Earth or Sun are unimportant,
• radiation appears to be quite isotropic.
• This is expected at all but the very highest
  energies
• since galactic magnetic fields would destroy any
  initial anisotropy.

8
Geomagnetic and solar effects

• Primary radiation, charged particleslt10 GeV, show
  directional and time effects.
• Charged primaries affected
• -by Earths magnetic field, like magnetic
  dipole.
• -by modulation in time due to solar wind (11y
  solar cycle).
• Geomagnetic effects -Axis of dipole is at an
  angle to axis of Earths rotation.
• -Geographical
  coordinates of poles varies slowly with time
• Calculation of actual orbits of particles
  incident on Earth as they spiral in dipole field
• Particle of charge ze, velocity v and momentum
  p mv
• travelling in circular equatorial path of radius
  r around a short dipole of moment M.
• Centrifugal and magnetic forces
  
• 
  ze B x v mv2/r
• Equatorial field due to dipole
  
• 
  B (µ0/4p) M/r3
• Radius of orbit(Stormer unit)
  
• 
  rS(µ0/4p) M ze/p1/2

9
Geomagnetic and solar effects

• Particle momentum that makes Earths radius rE
  equal to Stormer unit rS
• pc/z (µ0/4p) Mc e/(rE) 2 59.6 GeV
• Proton of smaller momentum cant reach Earth from
  eastern horizon at magnetic equator
• Stormer equation of motion obeyed by a particle
• b rsin?cos?cos2?/r
• r distance of particle from dipole centre
• ? geomagnetic latitude
• ? angle between v and its projection in
• meridian plane OAB co-moving with particle.
• b closest distance of approach to dipole axis
• by a tangent to particle trajectory.

10
Geomagnetic and solar effects

• Since sin? lt 1, restrictions on values of b, r,
  ?
• for allowed trajectories of particles reaching
  Earth.
• b rsin?cos?cos2?/r
• b ? 2 is critical in determining which momenta
  are cut off by the Earths field.
• Solving for r and with b2 for cut-off momentum
  at any ? and ? is
  
• r cos2?/1(1-sin?cos3?)
  1/2
• Using again the momentum equation
• pc/z (µ0/4p) Mc e/(rS)2
  (µ0/4p) Mc e/(rE) 2 59.6 GeV
• pc/z r 2 59.6 GeV
• since we are concerned with particles arriving at
  the Earth r rE/rS.
  
• pc/z 59.6 GeV cos2?/1(1-sin?cos3?)1/2
  2
• For particles incident from vertical ? 0 and r
  0.5 cos2?
• Cut-off momentum
• 
  (pc)min(? 0) 14.9 z cos4? GeV
  

11
Geomagnetic and solar effects

• Vertical(? 0) cut-off momentum pc/z
• In Europe ? 50 N
  1.1 GeV
• At magnetic equator ? 0
  14.9 GeV
• pc/z 59.6 GeV cos2?/1(1-sin?cos3?)1/2 2
• For ?0(at magnetic equator)
• For particles from Eastern horizon (sin?1)
  59.6 GeV
• For particles from Western horizon (sin?-1)
  10.2 GeV 59.6/(1 v2)2
• East-west effect
• at all latitudes, more (positively charged)
  particles arrive from West than from East,
• because of lower momentum cut-off.
• The effect arises essentially because all
  positively charged particles are deflected in a
• clockwise spiral, as viewed from above the N
  pole.

12
Acceleration of cosmic rays

• How do cosmic rays obtain their colossal
  energies, up to 1020eV?
• Energy density in cosmic rays, coupled with their
  lifetime in the galaxy,
• required a power supply similar to the rate of
  energy generation in supernova shells.
• Total power requirement to accelerate cosmic rays
  in the disc
• WCR ?E p R2
  D/t 2 x 1041 J/y
• Average energy density ?E of cosmic rays in the
  galaxy 1eV/cm3.
• Our own galaxy R 15kpc and thickness D
  0.2kpc.
• t 3 x 106 y average age of a cosmic ray
  particle in the galaxy,
• before it diffuses out or is depleted and lost in
  interactions with interstellar gas.
• Type II supernova typically ejects a shell of
  material of about 10 MO(2 x 1031 kg),
• with velocity 107m/s into interstellar medium
  approximately 1/century in our galaxy.
• Power output per galaxy WSN 5 x 1042 J/y.
  
• Even galactic supernova rate is uncertain, an
  efficiency for shockwave to transmit E to
• cosmic rays of few is enough to account for
  total E in cosmic ray beam.

13
Acceleration of cosmic rays

• At present time, is not at all clear which
  mechanisms contribute predominantly to the
• acceleration of cosmic-ray particles.
• In contrast, extremely energetic cosmic rays are
  predominantly accelerated in
• pulsars, binaries, or in jets emitted from black
  holes or active galactic nuclei.
• For shock acceleration in supernova explosions
• shape of E spectrum of cosmic-rays can be derived
  from acceleration mechanism.

14
Acceleration of cosmic rays

• In each cycle of acceleration at shock-front,
  particle gets energy increment ?E aE.
• After n cycles
• In terms of final energy E, of acceleration
  cycles
• At each stage of the acceleration, particle can
  escape further cycles.
• P probability that particle stays for further
  acceleration.
• After n cycles, of particles remaining for
  further acceleration
• N0 initial number of particles. Substituting for
  n
• s-lnP/ln(la). N number of particles with n or
  more cycles, with energy gt E.
• Differential energy spectrum will follow power
  law dependence

15
Acceleration of cosmic rays

• How do we account for the form of the energy
  spectrum?
  
• Differential energy spectrum will follow power
  law dependence
• For shock-wave acceleration s 1.1
• Differential spectrum index -2.1, compared with
  observed -2.7.
• Steeper observed spectrum could be accounted for
  
• if escape probability (1-P) was E dependent.
• Shock-wave acceleration from supernovae shells
  appears capable accounting for E of
• cosmic ray nuclei of charge Ze up to about 100Z
  TeV (l014Z eV), but not beyond this.
• Other, largely unknown, aceleration mechanisms
  must be invoked
• for very highest energy cosmic rays.

16
Secondary cosmic radiation

• Primary particles will produce secondaries
• Atmosphere as target in accelerator beam.
• -Radiation length for ? and e- 36.66 g/cm2.
• ?tmosphere corresponds to 27 radiation lengths.
• -Interaction length for hadrons 90.0g/cm2.
• ?tmosphere corresponds to 11 interaction lengths.
  
• Compared with total atmospheric depth1030 gm/cm2
• Most commonly produced particles p p- p0.
• Charged p decay to µ and ?
• Neutral p initiate via their decay (p0 ?? ? ? )
  electromagnetic cascades,

17
Decay vs. interaction of p (Proper lifetime t
26 ns, mc2 0.139 GeV.

• Decay Mean free path before decay ?
  ?ct (? E/mc2, time dilation factor).
• For 1 GeV p ? 55
  m.
• Interaction Depends on how much material has
  the atmosphere.
• Depth x (gm/cm2)
• x X exp(-h/H)
  X1030g/cm2 H 6.5 km
• In interval ?h(?55m)
  0.01H, depth change only by 1.
• At GeV almost all charged p decay in flight
  (rather than interact).
• __________________________________________________
  __________________________________________________
  __________________________________________________
  ______________________
• Decay At TeV, p decay probability
  100 sec?/EpGeV. ? zenith angle
• For E100GeV, mean free
  path before decay ? ?ct 5.5 km
• Interaccion Depth x (gm/cm2)
• x X exp(-h/H)
  X1030g/cm2 H 6.5 km
• Nuclear absorption
  important for charged p with ? H or Egt100GeV
• At TeV majority of p undergo nuclear interaction
  before they have a chance to decay

18
Secondary cosmic radiation

• Leptonic decays of p produce penetrating µ and ?
• µ can also decay and contribute
• -via their decay e- to soft component
• -via their decay ? to ? component
• ?nergy loss of relativistic µ not decaying in
  atmosphere is low .
• They constitute with 80 of all charged
  particles,
• the largest fraction of secondary particles at
  sea level.
• Some secondary mesons and baryons can survive to
  sea level.
• Most of low-? charged hadrons at sea level are
  locally produced.
• Total fraction of hadrons at ground level is very
  small.

19
Secondary cosmic radiation

• Daughter µ are also unstable Decay proper
  lifetime of t 2200 ns. mµ0.105 GeV.
• For 1 GeV µ Mean decay length ? ?ct (?
  E/mc2) (1/0.105) c 2200 6.6 km,
• equal to H of
  atmosphere
• x X exp(-h/H)
  H 6.5 km
• µ of lt1 GeV will decay in flight in the
  atmosphere.
• No competition with nuclear interaction since µ
  do not have strong interactions.
• __________________________________________________
  __________________________________________________
  __________________________________________________
  ________________________
• For 3 GeV µ Mean decay length ?20 km,
• typical distance from
  point of production to sea-level.
• Ionization E loss
  2MeV/gm/cm2 of air traversed(very low).
• Hard component of cosmic radiation
• µ of gt3GeV can get through entire atmosphere
  without decaying or arriving to rest.
• Higher E µ can reach deep underground.

20
Secondary cosmic radiation

• Intensity of p, e-, and µ of all E as function
  of altitude in atmosphere.
• ?bsorption of p exponential function.
• p are mostly primaries and interact depending on
  
• quantity of material in atmosphere.
• e- and e produced through p0 decay
• with subsequent pair production
• reach maximum intensity at altitude of 15 km
• soon after are relatively quickly absorbed.
• While flux of µ is attenuated relatively weakly.
• µ produced by charged p and Kaons.
• Mostly µ lt3GeV can decay in atmosphere,
• ? is smaller than available d in atmosphere
  before reaching the Earth.
• When µ is in inclined paths, decay probability
  increase(larger d in atmosphere).

21
Secondary cosmic radiation

• Because of steepness of E spectra, particle
  intensities are dominated by low-E particles
• Low-E particles are mostly of secondary origin.
• Momentum cut for geomagnetic field for lt GeV.
• If only particles with Egt1GeV are counted
• -Primary nucleons (pneutrons) with initial high
  E
• dominate over all particles down to altitudes of
  9km,
• where µ take over(secondary particle).
• At high altitude mostly primaries(mostly p).
• Only for p,same graph as for all E with
  different scale.
• -µ gt1GeV less absorved hard component
• -Low interaction probability of ?
• ? are practically not at all absorbed in
  atmosphere.
• Flux increases monotonically because additional
• ? are permanently produced by particle decays.
• E spectrum of primary particles is steep, E
  distribution of secondaries also reflect it.

22
?e and ?µ at sea level

• Besides charged particles, ?e ?µ are produced
• in kaon, p and µ decays atmospheric neutrinos
• ?ig background for neutrino astronomy.
• Propagation of atmospheric ? has provided new
  insights
• for elementary particle physics, such as ?
  oscillations.
• Verticalhorizontal ? spectra similar as for µ
  spectra.
• Parent of ? are dominantly p and kaons,
• their decay probability is increased compared to
• interaction probability at inclined directions,
• horizontal ? spectra are also harder in
  comparison
• to the spectra from vertical directions.
• ?µ would appear to dominate

23
?e and ?µ at sea level

• ?µ would appear to dominate, since
• are strongly suppressed.
• p and kaons almost exclusively produce ?µ.
• In µ decay equal numbers of ?e and ?µ are
  produced.
• At high E also semileptonic decays of charmed
  mesons constitute a source for ?.
• Based on these classical considerations the
  integral ? spectra yield a ?-flavour ratio of
• This ratio is modified by propagation effects
  like ? oscillations.

24
Secondary cosmic radiation

• Apart from longitudinal development,
• ?? hadronic cascades spread out laterally in
  atmosphere.
• Lateral size of EM cascade caused by multiple
  scattering of e- and e, in
• hadronic cascades caused by transverse momenta at
  production of secondary particles
• Comparison of shower
• development of 100TeV
• ? and 100TeV protons
• in atmosphere.
• Transverse momenta of
• secondary particles fan
• out the hadron cascade.

25
Secondary cosmic radiation

• Comparison of shower development of 1TeV ? and
  p over Chicago ground.
• Comparison of simulations of p and iron shower,
  simulations of ? and p.

26
Development of electromagnetic showers

• Longitudinal development of EM shower
• e- of initial E0 traversing a medium.
• 1. radiation length e- radiates ?, of energy
  E0/2.
• ?ext radiation length
• ? converts to e--e pair, each with E0/4.
• ?riginal e- radiates a further ?, also of E0/4.
• ?fter 2 radiation lengths 1 ?, 2 e-, 1 e, each
  of E0/4.
• In this way, after t radiation lengths e-, e, ?
  in equal numbers,
• each with E(t) E0/2t.
• We neglected ionization losses.
• Cascade multiplication process continues until
  particle ? falls to E EC, critical E,
• when ionization loss suddenly becomes dominant
  and
• no further radiation or pair conversion processes
  are possible.
• Cascade reaches a maximum and then ceases
  abruptly.

27
Development of electromagnetic showers

• -Shower maximum is at depth t t(max)
  ln(E0/Ec)/ln 2
• It increases logarithmically with primary energy
  E0.
• - of particles at shower maximum is N(max)
  2t(max) E0/Ec
• Proportional to primary energy.
• -Number of shower particlesgt E equal to number
  created at depths less than t(E)
• Differential energy spectrum of particles
• -Total track-length integral (of charged
  particles) in radiation lengths
• From E conservation since ionization loss per
  particle is Ec per radiation length,
• essentially all incident energy is finally
  dissipated as ionization E loss.
• Track-length integral gives a measurement of
  primary E.

28
Development of electromagnetic showers

• Effects of both radiation loss ionization loss
  are present throughout shower process.
• Actual shower consists of
• -an initial exponential rise,
• -a broad maximum and
• -a gradual decline thereafter.
• Simple model reproduces many of essential
  quantitative features of actual EM cascades.
• Model treated shower as one-dimensional.
• Actual showers spread out laterally, due mostly
  to Coulomb scattering of e-.
• Lateral spread of a shower

29
Extensive air showers

• EAS Cascades initiated by energetic primary
  particles which develop in atmosphere.
• If primary is high energy p rather than e-
  nuclear cascade develop through atmosphere.
• Nuclear interaction length in air 100 gm/cm2.
• p(or heavier nucleus) generates mesons,
• and they can in turn generate further particles
  in subsequent collisions.
• While in EM shower, e- lose bulk of their E in a
  radiation length,
• nucleons can penetrate through several
  interaction lengths,
• losing only fraction of their E in each
  encounter, to nucleons as well as mesons.
• In air, nuclear interaction length is 2.5 times
  radiation length
• cascades initiated by nucleons are much more
  penetrating than EM cascades.

30
Extensive air showers

• EAS has ??, muonic, hadronic, and ? component
• ?ir shower develops shower core of energetic
  hadrons(nucleons), which inject
• permanently ? into more widely spread ??
  component (by fresh p0 production and
• decay and fresh ?? cascades) and other shower
  components.
• p0, which are produced in nuclear interactions,
• whose decay ? produce e- and e via pair
  production,
• continually fed e-,e and ? component.
• e-,e and ? initiate ?? cascades through
  alternating processes of
• pair production and bremsstrahlung.
• µ and ? components are formed by decay of charged
  p and kaons.

31
Extensive air showers

• Development of EM cascades is shown for various
  primary E.
• Particle intensity increases initially in
  parabolical fashion and decays exponentially
• after maximum of shower has been reached.
• Longitudinal profile of particle parameterized
• N(t) ta e-ßt
• t x/X0 shower depth in units of radiation
  length
• a and ß are free fit parameters.
• Position of shower maximum varies
• only logarithmically with primary E.
• Total of shower particles increases linearly
  with E.
• (used for E determination of primary particle).
• Earths atmosphere represents a combined hadronic
  and EM calorimeter.
• Critical E in air is Ec 84 MeV.

32
Extensive air showers(EAS)

• Atmosphere is a target of 11 interaction
  lengths 27 radiation lengths.
• Minimum E for primary particle to be reasonably
• measured at sea level via particles
• produced in air shower is 1014 eV.
• Longitudinal development of components of
• EAS in atmosphere for primary ?1015 eV
• Particle N, at sea level in its
• dependence on primary E E0
• N 10-10 E0eV .

33
Neutrinos

• Neutrino physics is largely an art of learning a
  great deal by observing nothing.
• Limits of classical astronomy observations in
  radio, infrared, optical, UV,X-ray, ?-ray
• have limit due to EM radiation is quickly
  absorbed in matter.
• One can only observe the surfaces of astronomical
  objects.
• Energetic particles from distant sources are
  attenuated via interactions with ? of CMB.
• Requirement for an optimal astronomy
• 1. Optimal astroparticles or radiation should not
  be influenced by magnetic fields.
• 2. The particles should not decay from source to
  Earth.
• 3. Particles and antiparticles should be
  different for distinguish origin from matter or
  antimatter source.
• 4. Particles must be penetrating so that one can
  look into the central part of the sources.
• 5. Particles should not be absorbed by
  interstellar or intergalactic dust or IR or CMB
  ?.
• These requirements are fulfilled by ? in an ideal
  way!
• Why ? astronomy has not been a major branch of
  astronomy?
• ? can escape from the center of the sources
  because its low interaction cross section.
• But also enormous difficulty to detect ? on
  Earth.

34
Neutrinos

• For solar ? several 100 keV, cross section for
  ?nucleon scattering
• s(?eN) 10-45 cm2/nucleon
• Interaction probability of ? with our planet
  Earth at central incidence is
• f s NA d ? 4 ? 10-12
• NAAvogadro number,
• d diameter of the Earth,
• ? average density of the Earth.
• From 71010?/cm2s radiated by Sun arriving at
  Earth only 1 at most is seen here.
• ? telescopes must have an enormous target mass,
  and long exposure times.
• For high E, interaction cross section rises with
  ? E.
• ? several 100 keV can be detected by
  radiochemical methods.
• For ?gt 5 MeV large-volume water Cherenkov
  counters are an attractive possibility.
• Neutrino astronomy is a very young branch of
  astroparticle physics.

35
High energy Neutrinos

• Is considered realistic that point source in our
  galaxy produces ? spectrum according to
• This leads to an integral flux of ?
• The interaction cross section of high-energy ?
  was measured at accelerators
• s(?µN) 6.7 10-39 E? GeV cm2/nucleon.
• For 100 TeV ? s 6.7 ? 10-34 cm2/nucleon.
• For target thickness of 1km, an interaction
  probability W per ?
• W NA s d 4 ? 10 -5
• For d 1 km, ?(ice) 1 g/cm3.
• Total interaction rate R obtained from integral ?
  flux F? , interaction probability W,
• effective collection area Aeff 1 km2, and a
  measurement time t .
• Event rate R F? W Aeff 250
  events/year.

36
High energy Neutrinos

• Assuming half a dozen sources in our galaxy,
  estimate would be 4 events/day.
• Besides point sources,
• events from diffuse ? background carries little
  astrophysical information.
• Excellent candidates within our galaxy supernova
  remnants of Crab Nebula and Vela,
• galactic center, and Cygnus X3.
• Extragalactic candidates could be Markarian
  galaxies Mrk421, Mrk 501, M87, quasars.
• ? astronomy was pioneered with Baikal telescope
  installed in lake Baikal in Siberia.
• The most advanced larger telescopes are
  AMANDA/ICECUBE in the South pole and
• NESTOR/ANTARES/???? detector in the Mediterranean
  coasts of Greece, France
• and Italy.
• Real ? astronomy will require larger detectors.
• Auger experiment could also detect high energy ?
  events.

37
Gamma Astronomy

• An important, so far unsolved problem of
  astroparticle physics origin of cosmic rays.
• Investigations of charged primary cosmic rays are
  essentially unable to answer this
• because charged particles have to pass through
  extended irregular magnetic fields on
• their way from source to Earth.
• This causes them to be deflected in an
  uncontrolled fashion, forgetting their origin.
• Particle astronomy with charged particles is only
  possible at extremely high E when are
• no significantly affected by cosmic magnetic
  fields.
• Would require E gt1019 eV with flux of primary
  particles extremely low.
• Whatever the sources of cosmic rays are, they
  will also emit energetic penetrating ? rays
• not deflected by intergalactic or stellar
  magnetic fields and so point back to sources.
• However ? rays from distant sources might be
  subject to time dispersions.
• Astronomical objects in the line of sight of
  these sources can distort their trajectory by
• gravitational lensing making them look blurred
  and causing time-of-flight dispersions in
• arrival time also for electromagnetic radiation.

38
Extensive air showers

• Even EAS initiated by primary particles with
  Elt100TeV doesnt reach sea level,
• it can be recorded via Cherenkov light emitted by
  shower particles.
• As relativistic particles traverse the
  atmosphere,
• part of E loss appears as a coherent wavefront of
  Cerenkov radiation.
• This radiation is mostly in ultraviolet or blue
  region of the spectrum.
• Huyghens construction for emission of Cerenkov
  light by a relativistic particle
• Cos ? (ct/n)/(ßct)1/(ßn)
  ßgt 1/n
• Refractive index n of air at ground en-1
  0.0003.
• Threshold E for e- mc2/(1-ß2)1/2
  mc2/(2e)1/221MeV.
• For µ 4.3GeV (mµ 106 MeV).
• Most components of EAS have much greater E,
• so they will produce abundant Cerenkov light.
• Relativistic particle above threshold 10000?/km
  of path.
• Can be detected by large mirrors pointing to
  photomultipliersWhipple Observatory

39
Extensive air showers

• Cerenkov light emitted in narrow cone of
  ?(2e)1/2(1.4 at ground, although Coulomb
• scattering of e- will considerably broaden this),
  so that light appears in restricted radius
• 100m around shower axis, axis must be fairly
  close to mirror system to be recorded.
• Ionizing particles can excite fluorescence from
  N2 in the atmosphere
• 5000?/km of track in the blue wavelength region.
• Fluorescent light is emitted isotropicallydistant
  showers several km away, not aimed
• towards mirror/photomultiplier system, can be
  detected.
• Sensitivity to highest E events greatly
  increased.
• Whipple Observatory 1.system employing these
• techniques with large mirror array.
• 2 spatially separated arrays of mirrors for
  stereo
• images to reconstruct shower profile.
• Distinguish showers initiated by primary ?
• develop early, contained in upper atmosphere,
• from those by nucleonsdevelop slowly,penetrating.
  
• Poor duty cycle cloudless, moonless nights.

40
?-rays

• Air Cerenkov telescopes interesting type of
  ?-ray detector.
• ?ypical detector must be flown with balloon or
  satellite above
• atmosphere to avoid absorption of ?-ray, Air
  Cerenkov telescope
• nullifies problem by making atmosphere part of
  detector!
• ?-rays interacting in atmosphere create an air
  shower.
• Depending on ? of initial cosmic ?-ray, may be
  thousands of
• electrons/positrons in cascade emitting Cerenkov
  radiation.
• Light is pancake-like 200m in diameter but 1m
  thickness.
• HESS,MAGIC,...
• VERITAS (Very Energetic Radiation Imaging
  Telescope Array System), built by
• a collaboration headed by the Whipple
  Observatory, uses an array of seven 10m optical
• reflectors for gamma-ray astronomy in 50 GeV - 50
  TeV.
• Solar Tower Atmospheric Cerenkov Effect
  Experiment (STACEE) will use primary
• collection mirror large field of 220 solar
  heliostat mirrors at National Solar Thermal
• Test Facility of Sandia National Laboratories.
• Milagro built by collaboration headed by Los
  Alamos National Laboratory, use 5000m2
• pool of water covered with light-tight barrier as
  Cerenkov detector to study TeV ?-rays.

41
?-rays

• Observing the Universe in ?-rays allows us to
  examine things which are happening that
• cannot be seen with ordinary telescopes and which
  are very important in helping us to
• understand how matter and radiation interact with
  each other.
• This is especially true for understanding their
  interaction under extreme conditions
• where temperatures are hundreds of millions
  degrees, matter is very dense, or magnetic
• fields are very strong.
• Some specific targets include
• solar flares
• Gamma-ray Bursts
• Black Holes and Neutron Stars
• Supernovae
• Pulsars
• Active Galaxies Seyferts and Quasars

42
Extensive air showers

• At higher E various detection techniques.
• Classical techniquesampling of shower particles
  at sea level with 1m2 scintillators or
• water Cherenkov counters.Auger 75y ago extended
  array of detectors in coincidence.
• Shower detectable at sea-level for primary
  Egt1015eV, when maximum is near ground.
• At mountain altitudes, threshold 100 TeV.
• E assignment for primary not very precise.
• Shower develops in atmosphere which acts as
• calorimeter of 27 radiation lengths thickness.
• Information on shower is sampled in only one,
• the last layer of calorimeter and
• coverage of this layer is typically only 1.
• Direction of incidence of primary obtained from
• arrival times of shower particles in different
• sampling counters as shower front crosses array.
• Since shower front is quite well defined.

43
Extensive air showers

• More advantageous to measure total longitudinal
  development of cascade in atmosphere
• Apart from directional Cherenkov radiation shower
  particles emit isotropic scintillation
• light in atmosphere.
• For particles with Egt1017 eV, fluorescence light
  of N is sufficiently intense to be
• recorded at sea level in presence of diffuse
  background of starlight.
• Detector consists of a system of mirrors and
  photomultipliers, which view the sky.
• Air shower passing through atmosphere near
  activates only those photomultipliers
• whose field of view is hit.
• Fired photomultipliers allow to reconstruct
  longitudinal profile of air shower.

44
Extensive air showers

• Total recorded light intensity to determine
  shower E more precise E assignments.
• Big disadvantage compared to classical air-shower
  technique only clear moonless night
• Arrangement of mirrors and photomultipliers in
  Flys Eye setup of Utah.
• Auger experiment array of sampling detectors
  complemented by telescopes measuring
• fluoresce light produced in atmosphere.
• Much larger acceptances of these telescopes
• in orbit Air Watch.

45
Extensive air showers

• -?as been tried to observe air showers via EM
  radiation emitted in radio band.
• Is believed this radio signal is caused by
  shower e- deflected in Earths magnetic
• field creating synchrotron radiation.
• Because of strong background in all wavelength
  ranges,
• these attempts have not been particularly
  successful so far.
• -Possibility to detect large air showers via
  their µ content in underground experiments
• has been followed up in recent experiments.

46
Ultrahigh Energy cosmic rays

• -AGASA
• -Flys Eye
• -HiRes
• -Telescope Array
• Auger Observatory
• Biggest ultra high energy cosmic ray detector on
  Earth
• http//www.auger.org/
• http//www.auger.org.ar/