1. Ice-fishing for Cosmic Neutrinos Subhendu Rakshit TIFR, Mumbai

2. Goals of neutrino astronomy Astrophysics: To explore astrophysical objects like AGN or GRBs. Find out sources of high energy cosmic rays. Main aim.. Particle physics: To explore beyond standard model physics options which may affect neutrino nucleon cross-sections at high energy. Other possibilities… Appeared in US particle physics roadmap! First step: To determine the incoming neutrino flux

3. Astrophysical motivations Historically looking at the same astrophysical object at different wavelengths revealed many details regarding their internal mechanisms A 3-pronged approach involving conventional photon astronomy, cosmic ray astronomy and neutrino astronomy will yield better results

4. Conventional astronomy with photons Ranges from 10 4 cm radio-waves to 10 -14 cm high energy gamma rays Pros:  Photons are neutral particles. So they can point back to their sources  photons are easy to detect as they interact electromagnetically with charged particles Cons:  Due to the same reason they get absorbed by dust or get obstructed  Very high energy photons on its way interact with cosmic microwave background radiation and cannot reach us

5. Cosmic ray astronomy Very high energy cosmic rays (protons, heavy nuclei,..) do reach us from the sky It is difficult to produce such energetic particles in the laboratory It is puzzling where they are produced and how they get accelerated to such energies!! Although they can be detected on Earth, it is not possible to identify the sources as their paths get scrambled in magnetic fields  A serious disadvantage! Only very high energy(>10 10 GeV) cosmic rays point back to their sources

6. Neutrino astronomy The suspected sources of very high energy photons and cosmic rays are believed to be the sources of neutrinos as well Pros: Neutrinos being weakly interacting reaches Earth rather easily Cons: Due to the same reason it also interacts rarely with the detector material ⇒ Large detector size!! Successful neutrino astronomy with the sun and supernova. Now it is time to explore objects like Active Galactic Nuclei or Gamma Ray Bursts Impressive range for future neutrino telescopes: 10 2 GeV to 10 12 GeV!

7. GeV TeV PeV EeV 1 PeV = 10 6 GeV 1 EeV = 10 9 GeV Underwater / ice Air shower Underground Neutrino detectors

8. Why a Km 3 detector? Estimations of the expected amount of UHE neutrinos can be made from the observed flux of cosmic rays at high energies. This limits the size of the detector However such estimations are quite difficult as many assumptions go in There can be hidden sources of neutrinos!! So the neutrino flux can always be higher!

10.  IceCube o1KM^3o1KM^3 A Km 3 detector PMTs detect Cherenkov light emitted by charged particles created by neutrino interactions The Cherenkov cone needs to be reconstructed to determine the energy and direction of the muon

11. - The predecessor of IceCube Used for calibration, background rejection and air- shower physics

12. IceCube is optimised for detection of muon neutrinos above 1 TeV as: We get better signal to noise ratio Neutrino cross-section and muon range increases with energy. Larger the muon range, the larger is the effective detection volume The mean angle between muon and neutrino decreases with energy like 1/√E, with a pointing accuracy of about 1 ◦ at 1 TeV The energy loss of muons increases with energy. For energies above 1 TeV, this allows us to estimate the muon energy from the larger light emission along the track

13. IceCube Cosmic rays produce muons in our atmosphere, which can fake a neutrino-induced muon signal  background So we use the Earth to filter them out! Upto PeV neutrinos can cross the Earth to reach IceCube For high energy neutrinos Earth becomes opaque as the probability that the neutrinos will interact becomes higher with energy So very high energy neutrinos can reach Icecube only from the sky or from horizontal directions! Detection strategy

14. Sources of neutrinos Signal: The neutrinos from astrophysical sources: AGN or GRBs for example Background: Atmospheric neutrinos. They are produced from cosmic ray interactions with the atmosphere  A guaranteed flux well measured in AMANDA. Agrees with expectations. As the ATM flux falls rather rapidly( ∝ E -3 ) with energy, at higher energy we can observe the ‘signal’ neutrinos from AGN or GRBs free of these background neutrinos

15. Neutrino spectra Note: At higher energies the flux is smaller. But higher energy neutrinos also have higher cross- section. So detection probability is also higher!

16. Another background Cosmogenic or GZK neutrinos: UHE cosmic ray protons interact with CMBR photons to produce these neutrinos via charged pion decay However at IceCube the rate would be quite small

17. Eliminating backgrounds Energy cuts Directional cuts Directional signals Temporal considerations

19. Production at astrophysical sources: Initial flavour ratio Propagation through space: Massive neutrinos undergo quantum mechanical oscillations. So neutrinos reach Earth with a flavour ratio Propagation through the Earth: Neutrinos while propagating may interact with the Earth. CC or NC interactions. τ propagation is more elaborate: τ →τ→ τ →τ... Detection at IceCube: Muon neutrinos produce muons via CC interactions. All neutrinos produce showers through NC interactions. A CC interaction by a τ may produce spectacular signatures!

20. Production at astrophysical sources: A proton gets accelerated and hits another proton or a photon. They produce neutron, π + and π 0.Their decay produces cosmic rays, neutrinos and photons respectively p +  → π + + n p +  → π 0 + p

21. For massive neutrinos flavour and mass eigenstates are different. This implies that a neutrino of a given flavour can change its flavour after propagating for sometime! For example: µ ↔ e Neutrino oscillation At time t=0, we produce a e After sometime t, the mass eigenstates evolve differently So the probability of detecting another flavour is nonzero Propagation through space:

22. Now remember the initial flavour ratio at source was Recent neutrino experiments have established that neutrino flavour states µ and τ mix maximally Hence it is of no wonder that after traversing a long distance these two states will arrive at equal proportions Note that although there were no tau neutrinos at the source, we receive them on Earth! At source On Earth

23. While traversing through the Earth, neutrinos can undergo  a charged current(CC) interaction with matter. The neutrino disappears producing e or mu or tau. The dominant effect  or a neutral current interaction(NC) with matter. The neutrino produces another neutrino of same flavour with lower energy As a consequence, the number of neutrinos decrease as they propagate through the Earth. This depends on the energy of the neutrino. Higher energy neutrinos get absorbed more, their mean free path is smaller Propagation through the Earth:

25. µ detection Muons range: few Kms at TeV and tens of Km at EeV The geometry of the lightpool surrounding the muon track is a Km-long cone with gradually decreasing radius Initial size of the cone for a 100TeV muon is 130m. At the end of its range it reduces to 10m. The kinematic angle of µ wrt the neutrino is µ is 1 ◦ /√(E /1TeV) and the reconstruction error on the muon direction is on the order of 1 ◦ Better energy determination for contained events. More contained events at lower energy

26. ~ Km long muon tracks from µ ~ 10m long cascades from e, τ

27. e detection In a CC interaction, a e deposits 0.5-0.8% of their energy in an EM shower initiated by the electron. Then a shower initiated by the fragments of the target The Cherenkov light generated by shower particles spreads over a vol of radius 130m at 10TeV and 460m at 10EeV. Radius grows by ~50m per decade in energy Energy measurement is good. The shower energy underestimates the neutrino energy by a factor ~3 at 1 TeV to ~4 at 1 EeV Angle determination poor! Elongated in the direction of e so that the direction can be reconstructed but precise to ~10 ◦

28. The propagation mechanism of a tau neutrino is different, as tau may decay during propagation As a result the tau neutrino never disappears. For each incoming τ another τ of lower energy reaches the detector The Earth effectively remains transparent even for high energy tau neutrinos Tau decays produce secondary flux of e and µ τ τ τ τ τ detection

29. Double bang events: CC interaction of τ followed by tau decay Lollipop events: second of the two double bang showers with reconstructed tau track Inverted lollipop events: first of the two double bang showers with reconstructed tau track. Often confused with a hadronic event in which a ~100GeV muon is produced! For E τ < 10 6 GeV, in double bang events showers are indistinguishable. For E τ ~ 10 6 GeV, tau range is a few hundred meters and the showers can be separated. For 10 7 GeV < E τ < 10 7.5 GeV, the tau decay length is comparable to the instrumented detector vol.  lollipop E τ > 10 7.5 GeV tau tracks can be confusing

30. Propagation equation of µ

31. Propagation equations of τ

34. Including energy loss Without energy loss

35. Characteristic bump Rakshit, Reya, PRD74,103006(2006)

37. Expected muon event rate per year at IceCube µ induced µ + τ induced

38. Imprinted Earth’s matter profile

40. Production at astrophysical sources: Initial flavour ratio ? Propagation through space: Massive neutrinos undergo quantum mechanical oscillations. So neutrinos reach Earth with a flavour ratio ?? Propagation through the Earth: Neutrinos while propagating may interact with the Earth. CC or NC interactions. τ propagation is more elaborate: τ →τ→ τ →τ... Detection at IceCube: Muon neutrinos produce muons via CC interactions. All neutrinos produce showers through NC interactions. A CC interaction by a τ may produce spectacular signatures! N xsection sensitive

41. Detection of atm µ s will enable us to probe CPTV, LIV,VEP which change the standard 1/E energy dependence of osc length. Due to high threshold of IceCube, osc of these high energy atm neutrinos is less N xsection can get enhanced in XtraDim models N xsection can get reduced at high energies in color glass condensate models Visible changes in muon rates, shower rates For xtradim upgoing neutrinos get absorbed at some energy and also downgoing for higher energies For lower N xsection models angular dependence and energy dependence for upgoing events are more important

42. Crude neutrino flux determination from up/down events OK for fixed power flux, but otherwise contained muon events are better. But poorer statistics Auger is better for UHE neutrinos. New physics effects will be more dramatic IceCube can probe neutrino spectrum better as Xsection uncertainties are only at high energies where the flux is smaller Flavour ratio determination possible at IceCube as different flavours have distinctive signatures.

43. Other possibilities DM detection: Neutrinos from solar core SUSY search: look for charged sleptons RPV, Leptoquarks Part of supernova early detection system! New physics interactions at the detector New physics during propagation

44. Summary UHE neutrinos: particle physics opportunities for the future IceCube is a discovery expt. Determining neutrino spectrum independent of new physics poses a challenge Even crude measurements at IceCube may provide some clue about drastically different new physics scenarios at high energies Some success with IceCube will lead to bigger detectors At present we just need to detect an UHE neutrino event at IceCube!

45. Particle physics motivations LHC CM energy E CM = 14 TeV ⇒ LHC: E =10 8 GeV Tevatron: E =10 6 GeV Here we talk about neutrino flux of 10 12 GeV! ⇒ E CM = 14 ×100 TeV

46. N cross-sections We need PDF’s for x 10 8 GeV Several options but not much discrepancy! GRV and CTEQ cross-sections differ at the most by 20%

47. Horizontal μ creating a detectable μ track For downgoing μ e shower(CC+NC) τ lollipop τ double bang Beacom et al, PRD 68,093005(2003)