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Neutrino astronomy and telescopes

Neutrino astronomy and telescopes. Crab nebula. Cen A. Teresa Montaruli, Assistant Professor, Chamberlin Hall, room 5287, tmontaruli@icecube.wisc.edu. Neutrinos and their properties. Neutrino astronomy and connections to Cosmic rays and gamma-astronomy.

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Neutrino astronomy and telescopes

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  1. Neutrino astronomy and telescopes Crab nebula Cen A Teresa Montaruli, Assistant Professor, Chamberlin Hall, room 5287, tmontaruli@icecube.wisc.edu

  2. 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 Overview Teresa Montaruli, 5 - 7 Apr. 2005

  3. Some neutrino hystory • 4 Dec 1930:W. Pauli pioneering hypothesis on neutrino existence as a “desperate remedy” to explain the continuous b-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…. • 1933 E. Fermi :b-decaytheory Week interactions: GF << a of electromagnetic interactions • 1956 Cowan and Reines : first detection of reactor neutrinos • by simultaneous detection of 2g‘s from e+ pair annihilation and • neutron Teresa Montaruli, 5 - 7 Apr. 2005

  4. Astrophysical neutrinos: from the Sun Combined effect of nuclear fusion reactions Predicted fluxes from Standard Solar Model Uncertainty ~ 0.1% Pioneer experiment: 1966 R. Davis in Homestake Mine Radiochemical experiment: 615 tons of liquid perchloroethylene (C2Cl4), reaction ne + 37Cl -> e- + 37Ar, Eth=0.814 MeV, operated continuously since 1970 Observed event rate of 2.56±0.23 SNU (1 SNU =10-36 interactions per target atom per second)  Standard Solar Model prediction: 7.7+1.2-1.0 SNU Solar neutrino problem, now solved by oscillations Teresa Montaruli, 5 - 7 Apr. 2005

  5. SN1987A:99% of binding energy in ns in a core collapse SN Neutronization, ~10 ms 1051 erg • Thermalization: ~10 s 31053 erg Astrophysical neutrinos: from SN1987A http://www.nu.to.infn.it/Supernova_Neutrinos/#7 Teresa Montaruli, 5 - 7 Apr. 2005

  6. The challenge We learned: Weak interactions make neutrinos excellent probes of the universe but their detection is difficult ! Teresa Montaruli, 5 - 7 Apr. 2005

  7. Atmospheric ns ns from WIMP annhilation Cosmic ns Neutrino Fluxes Teresa Montaruli, 5 - 7 Apr. 2005

  8. Why neutrinos are interesting? • After photons (400 g/cm3) is the most abundant element from the Big Bang in the Universe (nn~3/11ng) • Open questions: mass? Majorana or Dirac? Leptons and quarks in Standard Model are Dirac particles: particles differ from antiparticles, 2 helicity states In the Standard Model the n is massless and neutral and only nL and nR. It is possible to extend the SM to have massive neutrinos and they may be Majorana particles (particle=antiparticle) if only nL ad nR exist The mass is a fundamental constant: needs to be measured!! Direct neutrino mass measurements ne< 3 eV ~ 3 x 10-9 mproton from b decay of 3H (Z,A)(Z+1,A) + e- + ne nµ< 0.17 MeV ~ 2 x 10-4 mproton from pmnm nt< 18.2 MeV ~ 2 x 10-2 mproton from t5pp0nt Neutrino mass =0? Teresa Montaruli, 5 - 7 Apr. 2005

  9. Neutrino properties: oscillations A n 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 Dm2~<E>/L depending on the experiment design Teresa Montaruli, 5 - 7 Apr. 2005

  10. Recent atmospheric neutrino experiments (Super-Kamiokande, MACRO, Soudan 2) have demonstrated that the nm deficit is due to nm nt oscillations with maximal mixing Dmatm2~ 2.5 · 10-3 eV2 sin22q23 ~ 1 Solar neutrino experiments: (Cherenkov detectors: Super-Kamiokande, SNO)+ KAMLAND: scintillator detector looking for ne from reactors at ~180 km average distance) deficit compatible with Dmsun2~ 7.1 · 10-5 eV2 sin22q12 ~ 0.82, could be due to nent or ne  nm Reactor neutrino experiments L~1 km (CHOOZ) constrain the q13 mixing (no disappearance) 3 2 1 SNO at Sudbury Mine The experimental scenario Teresa Montaruli, 5 - 7 Apr. 2005

  11. straight line propagation to point back to sources Photons:reprocessed in sources and absorbed by extragalactic backgrounds For Eg > 500 TeV do not survive journey from Galactic Centre Protons: directions scrambled by galactic and intergalactic magnetic fields (deflections <1° for E>50 EeV) Interaction length p + gCMBp + n lgp =( nCMBs ) -1~10 Mpc Neutrons: decay gct  E/mn ct ~10kpc for E~EeV Astronomy with particles q d Rgyro evB = mv2/Rgyro eB = p/Rgyro  1/Rgyro = B/E Teresa Montaruli, 5 - 7 Apr. 2005

  12. W49B g n SN 0540-69.3 p p + gp + n pgee Crab 3C279 Mrk421 g+IRe+e- Cas A Local Group g+radio E0102-72.3 g+MW Gal Cen <100 Mpc 1-109 TeV Messengers from the Universe 1 pc ~ 3 ly ~ 1018 cm 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 Teresa Montaruli, 5 - 7 Apr. 2005

  13. 106 eV to ~1020 eV ~E-2.7 ~E-3.1 knee Balloons satellites EAS Ankle: 1 km-2 century-1 ~E-2.7 After T. Gaisser, ICHEP02 The CR spectrum SN provide right power for galactic CRs up to the knee: CR energy density: rE~ 1 eV/cm3 ~ B2galactic/ 8p Needed power: rE / tesc~10-26 erg/cm3s with galactic escape time tesc ~ 3 x 106 yrs SN power: 1051 erg/SN + ~3 SN per century in disk ~ 10-25 erg/cm3s  10% of kinetic energy in proton and nuclei acceleration Teresa Montaruli, 5 - 7 Apr. 2005

  14. Hillas Plot R = acceleration site dimensions CR acceleration at sources The accelerator size must be larger than Rgyro energy losses in sources neglected Teresa Montaruli, 5 - 7 Apr. 2005

  15. Kascade - QGSJET The knee • What is the origin of the knee? • Acceleration cutoff Emax~ZBL~Zx100TeV, change in acceleration process? • Confinement in the galactic magnetic field: rigidity dependent cut-off • Change in interaction properties (eg. onset of channel where energy goes into unseen particles) 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 Teresa Montaruli, 5 - 7 Apr. 2005

  16. AGASA: 111 scintillators + 27 m detectors The ankle: the EHE region • 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”? Ankle: E-2.7 at E~1019 eV could suggest a new light population Protons are favored by all experiments. Fe frac. (@90% CL): < 35% (1019 –1019.5 eV), < 76% (E>1019.5eV) Gamma-ray fraction upper limits(@90%CL) 34% (>1019eV)(g/p<0.45) 56% (>1019.5eV)(g/p<1.27) Teresa Montaruli, 5 - 7 Apr. 2005

  17. Controvariant pa = (E/c,px,pypz)=(E/c,p) Scalar product of 2 4-vectors qaka = (EqEk/c2-pqxpkx-pqypqy-pqzpqz) Square p2 = (E/c)2 - p2 = m2 = constant Relativity: 4-vectors Covariant pa = (E/c,-px,-py-pz) Transform from one coordinate system to another moving with speed v in the x direction (Lorentz transformation) p’x = g(px – bE/c) p’y = py p’z = pz E’ = g(E – bpxc) Also: In general: (E*,p*) in a frame moving at velocity bf: ||=parallel to direction of motion T=transverse Teresa Montaruli, 5 - 7 Apr. 2005

  18. Energy of projectile to produce particles in the final state at rest mp mt at rest Reaction Thresholds mt,pt mt,pp s = Ecm2 c= 1 True in any reference system In the lab Teresa Montaruli, 5 - 7 Apr. 2005

  19. [Greisen 66; Zatsepin & Kuzmin66] 2.73 K Threshold for GZK cut-off Threshold pr p-gDpN in frame where p is at rest Energy of CMB photons: =3kBT effective energy for Planck spectrum And their energy in the proton rest frame is • gp= 2· 1011 and the threshold energy of the proton is then Ep = gp mp = 2 ·1020 eV Teresa Montaruli, 5 - 7 Apr. 2005 Integrating over Planck spectrum Ep,th~ 5 ·1019 eV

  20. [Greisen 66; Zatsepin & Kuzmin66] GZK cut-off? AGASA: 11 events, expects 2 @ E> 1020 eV 4s from GZK model from uniform distribution of sources Hires (fluorescence technique) compatible at 2s Uncertainties on E ~30% Not enough statistics to solve the controversy AGASA anisotropies: E>4 ·1019 eV • Air fluorescence detectors • HiRes 1 - 21 mirrors • HiRes 2 - 42 mirrors • Dugway (Utah) Teresa Montaruli, 5 - 7 Apr. 2005

  21. Anistropies Galaxy cannot contain EHECR: at 1019 eV Larmour radius of CR p comparable to Galaxy scale AGASA: E>4 ·1019 eVno evidence of anisotropies due to galactic disc butlarge 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/9505093 Teresa Montaruli, 5 - 7 Apr. 2005

  22. Berezinsky et al, 1985 Gaisser, Stanev, 1985 Neutrino production: bottom up Beam-dump model:p0 g-astronomy p± n-astronomy Neglecting g absorption (uncertain) n g Targets: p or ambient g Teresa Montaruli, 5 - 7 Apr. 2005

  23. From photon fluxes to n predictions:pp K = 1 pp 2 photons with 2nm and 1 ne with K = 1 since energy in photons matches that in nms 2nms with Ep/12 for each g Ep/6 Minimum proton energy fixed by threshold for p production (G =E/m is the Lorentz factor of the p jet respect to the observer) The energy imported by a n in p decay is ¼ Ep Teresa Montaruli, 5 - 7 Apr. 2005 Exercises!

  24. From photon fluxes to n predictions: pg K = 4 pg BR = 2/3 BR = 1/3 1) 2gs with 2/3× Eg = 2/3 ·0.1Ep K = 4 2) 2nms with 1/3× En = 1/3 0.1·Ep /2 Teresa Montaruli, 5 - 7 Apr. 2005

  25. Magnetic Cloud Magnetic inhomogenities 2nd order Fermi acceleration (1st 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 v Head-on Following This results in a net energy gain per collision of 2nd order in the velocity of the cloud Teresa Montaruli, 5 - 7 Apr. 2005

  26. CasA Supernova Remnant in X-rays Shock fronts Magnetic Inhomogenities v2>v1 John Hughes, Rutgers, NASA Interstellar medium Blast shock 1st order Fermi acceleration The 2nd order mechanism is a slow process. The 1st 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 1nd order in the velocity of the shock Equation of continuity: r1v1=r2v2 For ionized gas r1/r2 = 4  v2 =4 v1 upstream v1=u/4<v2 v2=u Shock front at rest: upstream gas flows into shock at v2 =u And leaves the shock with v1 = u/4 downstream Teresa Montaruli, 5 - 7 Apr. 2005

  27. Particles lost in each round trip on the shock Fermi mechanism and power laws E = bE0 = 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 = bkE0 and N = N0Pk = number of particles v1 velocity of gas leaving the shock v2 velocity of gas flowing into shock Naturally predicts CRs have steeper spectrum due to energy dependence of diffusion in the Galaxy Teresa Montaruli, 5 - 7 Apr. 2005

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