Neutrino astronomy and telescopes Teresa Montaruli, Assistant Professor, Chamberlin Hall, room 5287, Crab nebula Cen A

Teresa Montaruli, Apr. 2005Overview Neutrinos and their properties Neutrino astronomy and connections to Cosmic rays and gamma-astronomy Neutrino sources and neutrino production Neutrino telescopes Search Methods The Cherenkov technique and the photosensors Current experimental scenario

Teresa Montaruli, Apr Some neutrino hystory 4 Dec 1930: W. Pauli pioneering hypothesis on neutrino existence as a “desperate remedy” to explain the continuous  -decay energy spectrum Dear radioactive ladies and gentlemen, As the bearer of these lines, to whom I ask you to listen graciously, will explain more exactly, considering the ‘false’ statistics of N-14 and Li-6 nuclei, as well as the the desperate remedy……Unfortunately, I cannot personally appear in Tübingen, since I am indispensable here on account of a ball taking place in Zürich in the night from 6 to 7 of December… E. Fermi :  -decay theory Week interactions: G F <<  of electromagnetic interactions 1956 Cowan and Reines : first detection of reactor neutrinos by simultaneous detection of 2  ‘s from e + pair annihilation andneutron

Teresa Montaruli, Apr Astrophysical neutrinos: from the Sun Pioneer experiment: 1966 R. Davis in Homestake Mine Radiochemical experiment: 615 tons of liquid perchloroethylene (C 2 Cl 4 ), reaction e + 37 Cl -> e Ar, E th =0.814 MeV, operated continuously since 1970 Observed event rate of 2.56±0.23 SNU (1 SNU = interactions per target atom per second)  Standard Solar Model prediction: SNU  Solar neutrino problem, now solved by oscillations Combined effect of nuclear fusion reactions Predicted fluxes from Standard Solar Model Uncertainty ~ 0.1%

Teresa Montaruli, Apr Astrophysical neutrinos: from SN1987A SN1987A: SN1987A: 99% of binding energy in s in a core collapse SN Neutronization, ~10 ms erg Thermalization: ~10 s 3  erg

Teresa Montaruli, Apr The challenge We learned: Weak interactions make neutrinos excellent probes of the universe but their detection is difficult !

Teresa Montaruli, Apr Neutrino Fluxes Atmospheric s s from WIMP annhilation Cosmic s

Teresa Montaruli, Apr Why neutrinos are interesting? After photons (400  /cm 3 ) is the most abundant element from the Big Bang in the Universe (n ~3/11n  ) After photons (400  /cm 3 ) is the most abundant element from the Big Bang in the Universe (n ~3/11n  ) Open questions: mass? Majorana or Dirac? Open questions: mass? Majorana or Dirac? Leptons and quarks in Standard Model are Dirac particles: particles differ from antiparticles, 2 helicity states Leptons and quarks in Standard Model are Dirac particles: particles differ from antiparticles, 2 helicity states In the Standard Model the is massless and neutral and only L and R. In the Standard Model the is massless and neutral and only L and R. It is possible to extend the SM to have massive neutrinos and they may be Majorana particles (particle=antiparticle) if only L ad R exist It is possible to extend the SM to have massive neutrinos and they may be Majorana particles (particle=antiparticle) if only L ad R exist The mass is a fundamental constant: needs to be measured!! Direct neutrino mass measurements e < 3 eV ~ 3 x m proton from  decay of 3 H (Z,A)  (Z+1,A) + e - + e µ < 0.17 MeV ~ 2 x m proton from    < 18.2 MeV ~ 2 x m proton from   5    Neutrino mass =0?

Teresa Montaruli, Apr Neutrino properties: oscillations A created in a leptonic decay of defined flavor is a linear superposition of mass eigenstates Given a neutrino beam of a given momentum the various mass states have different energies and after a time t the probability that another flavor appears is where L=baseline For 2 flavor: Though oscillations are an indirect way of measuring the mass that requires many different experiments to reach an understanding of the difference of the square masses and of the flavors involved, they have the merit of being sensitive to very small masses  m 2 ~ /L depending on the experiment design

Teresa Montaruli, Apr The experimental scenario SNO at Sudbury Mine Recent atmospheric neutrino experiments (Super-Kamiokande, MACRO, Soudan 2) have demonstrated that the  deficit is due to    oscillations with maximal mixing  m atm 2 ~ 2.5 · eV 2 sin 2 2   ~ 1 Solar neutrino experiments: (Cherenkov detectors: Super-Kamiokande, SNO)+ KAMLAND: scintillator detector looking for e from reactors at ~180 km average distance) deficit compatible with  m sun 2 ~ 7.1 · eV 2 sin 2 2   ~ 0.82, could be due to e   or e   Reactor neutrino experiments L~1 km (CHOOZ) constrain the  13 mixing (no disappearance)321

Teresa Montaruli, Apr Astronomy with particles straight line propagation to point back to sources Photons: reprocessed in sources and absorbed by extragalactic backgrounds For E  > 500 TeV do not survive journey from Galactic Centre Protons: directions scrambled by galactic and intergalactic magnetic fields (deflections 50 EeV) p +  CMB   + n Interaction length p +  CMB   + n  p = ( n CMB  ) -1 ~ 10 Mpc  p = ( n CMB  ) -1 ~ 10 Mpc Neutrons: decay  ct  E/m n ct ~10kpc for E~EeV d R gyro evB = mv 2 /R gyro  eB = p/R gyro  1/R gyro = B/E 

Teresa Montaruli, Apr  p Messengers from the Universe Photons currently provide all information on the Universe but interact in sources and during propagation Neutrinos and gravitational waves have discovery potential because they open a new window on the universe p  ee p +   + n  +radio  +IR  e + e -  +MW Local Group 3C279 Mrk421 Gal Cen <100 Mpc TeV W49B SN Crab E Cas A 1 pc ~ 3 ly ~ cm

Teresa Montaruli, Apr ~E -3.1 ~E -2.7 After T. Gaisser, ICHEP eV to ~10 20 eV Balloonssatellites EAS knee Ankle: 1 km -2 century -1 The CR spectrum SN provide right power for galactic CRs up to the knee: CR energy density:  E ~ 1 eV/cm 3 ~ B 2 galactic / 8  Needed power:  E /  esc ~ erg/cm 3 s with galactic escape time  esc ~ 3 x 10 6 yrs SN power: erg/SN + ~3 SN per century in disk ~ erg/cm 3 s  10% of kinetic energy in proton and nuclei acceleration

Teresa Montaruli, Apr Hillas Plot R = acceleration site dimensions CR acceleration at sources energy losses in sources neglected The accelerator size must be larger than R gyro

Teresa Montaruli, Apr The knee Modest improvements in hadronic interaction models due to large uncertainties (different kinematic region than colliders) + stochastic nature of hadronic interactions  large fluctuations in EAS measurements Kascade - QGSJET What is the origin of the knee? 1.Acceleration cutoff 1.Acceleration cutoff E max ~ZBL~Zx100TeV, change in acceleration process? 2.Confinement in the galactic magnetic field: rigidity dependent cut-off 3.Change in interaction properties (eg. onset of channel where energy goes into unseen particles)

Teresa Montaruli, Apr The ankle: the EHE region Ankle: E -2.7 at E~10 19 eV could suggest a new light population Protons are favored by all experiments. What is the acceleration mechanism at these energies? Which are the sources? Are there extra-galactic components? Which particles do we observe? Is there the expected GZK “cutoff”? AGASA: 111 AGASA: 111 scintillators + 27  detectors Fe frac. CL): eV) Gamma-ray fraction upper 34% (>10 19 eV) (  /p eV) (  /p<1.27)

Teresa Montaruli, Apr Relativity: 4-vectors Covariant p  = (E/c,-p x,-p y- p z ) Scalar product of 2 4-vectors q  k  = (E q E k /c 2 -p qx p kx -p qy p qy -p qz p qz ) Square p 2 = (E/c) 2 - p 2 = m 2 = constant Also: Transform from one coordinate system to another moving with speed v in the x direction (Lorentz transformation) p’ x =  (p x –  E/c) p’ y = p y p’ z = p z E’ =  (E –  p x c) In general: (E*,p*) in a frame moving at velocity  f : ||=parallel to direction of motion T=transverse Controvariant p  = (E/c,p x,p y p z )=(E/c,p)

Teresa Montaruli, Apr Reaction Thresholds m t, p p m t,p t True in any reference system In the lab s = E cm 2 c= 1 mpmpmpmp m t at rest Energy of projectile to produce particles in the final state at rest

Teresa Montaruli, Apr Threshold for GZK cut-off [Greisen 66; Zatsepin & Kuzmin66] Integrating over Planck spectrum E p,th ~ 5 ·10 19 eV Threshold pr p-  N =3k B T effective energy for Planck spectrum 2.73 K Energy of CMB photons: in frame where p is at rest And their energy in the proton rest frame is  p = 2· and the threshold energy of the proton is then E p =  p m p = 2 ·10 20 eV

Teresa Montaruli, Apr GZK cut-off? AGASA: 11 events, expects E> eV 4  from GZK model from uniform distribution of sources Hires (fluorescence technique) compatible at 2  Uncertainties on E ~30% Not enough statistics to solve the controversy AGASA anisotropies: E>4 ·10 19 eV [Greisen 66; Zatsepin & Kuzmin66] Air fluorescence detectors HiRes mirrors HiRes mirrors Dugway (Utah)

Teresa Montaruli, Apr Anistropies Galaxy cannot contain EHECR: at eV Larmour radius of CR p comparable to Galaxy scale AGASA: E>4 ·10 19 eV no evidence of anisotropies due to galactic disc but large scale isotropy  EHECR are extra-galactic AGASA : 67 events cluster  1 triplet (chance prob <1%) + 9 doublets (expect 1.7 chance probability <0.1%) at small scale (<2.5˚) Not confirmed by HiRES Triplet close to super-galactic plane See also UHECR correlation with super-galactic plane astro-ph/

Teresa Montaruli, Apr Neutrino production: bottom up Neglecting  absorption (uncertain)    Targets: p or ambient  Beam-dump model:  0   -astronomy  ±  -astronomy Berezinsky et al, 1985 Gaisser, Stanev, 1985

Teresa Montaruli, Apr From photon fluxes to predictions:pp K = 1 pp 2 photons with 2  and 1 e with Minimum proton energy fixed by threshold for  production (  =E/m is the Lorentz factor of the p jet respect to the observer) The energy imported by a in  decay is ¼ E  K = 1 since energy in photons matches that in  s 2  s with E p /12 for each  E p /6 Exercises!

Teresa Montaruli, Apr From photon fluxes to predictions: p  K = 4 p  BR = 2/3 BR = 1/3 1) 2  s with 2/3× E  = 2/3 ·0.1E p 2) 2  s with 1/3× E = 1/3 0.1·E p /2 K = 4

Teresa Montaruli, Apr nd order Fermi acceleration (1 st version 1949) Magnetic clouds in interstellar medium moving at velocity V (that remains unchanged after the collision with a relativistc particle particle v~c) The probability of head-on encounters is slightly greater than following collisions V V Head-on Following This results in a net energy gain per collision of v 2 nd order in the velocity of the cloud Magnetic Cloud Magnetic inhomogenities

Teresa Montaruli, Apr CasA Supernova Remnant in X-rays John Hughes, Rutgers, NASA Shock fronts 1 st order Fermi acceleration The 2 nd order mechanism is a slow process. The 1 st order is more efficient since only head-on collisions in shock waves High energy particles upstream and downstream of the shock obtain a net energy gain when crossing the shock front in a round trip 1 nd order in the velocity of the shock Equation of continuity:  1 v 1 =  2 v 2 For ionized gas  1 /  2 = 4  v 2 =4 v 1 v 2 =u Shock front at rest: upstream gas flows into shock at v 2 =u And leaves the shock with v 1 = u/4 v 1 =u/4<v 2 Interstellar medium Blast shock Magnetic Inhomogenities v 2 >v 1 upstream downstream

Teresa Montaruli, Apr Fermi mechanism and power laws E =  E 0 = average energy of particle after collision P = probability of crossing shock again or that particle remains in acceleration region after a collision After k collisions: E =  k E 0 and N = N 0 P k = number of particles Naturally predicts CRs have steeper spectrum due to energy dependence of diffusion in the Galaxy Particles lost in each round trip on the shock v 1 velocity of gas leaving the shock v 2 velocity of gas flowing into shock