Neutrino Astronomy / Astro-physics

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1 Neutrino Astronomy / Astro-physics teresa.montaruli@unige.ch We covered: - Measured CR spectrum - Candidate Galactic and extra-galactic sources from energy balance between energy density of sources and injected in CRs - Non-thermal spectra (power laws): Fermi mechanisms. P esc = 1 T cycle esc E k = E 0 (1 + ) k 2 features: 1) The higher the energy of a particle (E0) the slowest the acceleration process (the smaller is ξ) 2) Emax will depend on the maximum number of cycles kmax = Tshock/Tcycle : E max = k max E A shock wave looses its accelerating efficiency when the ejected mass = mass of the swept up in ISM or ρ SN = ρ ISM ~1p/cm 3 => E max ZeβB Z x 300 TeV Then the Sedov phase (adiabatic shock deceleration) starts Niels Bohr Institute PhD School, 3 Aug. 2015

2 Neutrino Astronomy / Astro-physics Today - hadronic/leptonic processes - Neutrino beam dumps: pp and p-gamma hadronic processes - Flux of neutrinos and gammas from source injection proton spectrum - Effective areas Niels Bohr Institute PhD School, 1 Aug. 2015

3 MWL spectral emission distribution RXJ Psync/PIC = Umag/Urad in the Thompson regime energy flux E 2 d/de ~ νf(ν) kev synchrotron radiation from electrons de/dt Sy = k γ e 2 U mag ~ γ e 2 B 2 TeV Inv. Compton scattering by the same electrons de/dt IC = k γ e 2U rad photon energy [ev] In leptonic processes: the synchrotron luminosity is proportional to the B-field strength and Lorentz factor of electrons radiating. Inverse Compton: in the rest frame of an electron the photon energy is boosted by γe and after scattering, transforming in the lab again, the photon energy is boosted by γe 2. The IC luminosity is also proportional to the energy density of seed photons which can be the synchrotron photons. Synchrotron cooling is faster than IC if B > 3 ug. In hadronic processes, photons are produced via π 0 γγ.

4 Crab Nebula multiwavelength spectrum Aharonian et al Crab (2-10 kev) = erg cm -2 s -1 High B field (~160 µg) Synch. IC ~ 80% SSC ~ 20% IC FIR,mm, CMB Spectra extend to about 40 TeV

5 FERMI OBSERVATIONS OF RXJ Aharonian et al., Nature, 432, 75 (2004) ; Aharonian et al., A&A, 464, 235 (2007) `The dominance of leptonic processes in explaining the gamma-ray emission does not mean that no protons are accelerated in this SNR, but that the ambient density is too low to produce a significant hadronic gamma-ray signal. (arxiv: ). SNR maybe efficient accelerators during short periods when the shock speed is high and they transit from the free expansion (ejecta dominated phase) to the Sedov phase (adiabatic deceleration). Eg. Bell & Lucek, 2001

6 EVIDENCE FOR PION DECAY? Fermi identified various candidates for pionic gammas but yet did not verify that what we yet do not have a convincing model on CR acceleration Ackermann et al. (Fermi Collaboration), Science, 339, 807 (2013) Recent news: the First PeVatron from Galactic Centre with 10 yrs of HESS data Nature 531 (2016) 476 (

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8 Evolution of gamma-ray properties with age Gamma-ray properties evolve with SNR age and when the shock acceleration is very efficient photons may be not visible We need more source statistics! CTA Ahlers, Mertsch & Sarkar,PRD80:123017,2009

9 The general model for neutrino - gamma -CR sources ν and γ beam dumps accelerator (BH, n-star, shock ) is powered by large gravitational energy cosmic ray + gamma cosmic ray + neutrinos

10 pp and p-gamma processes about 2 orders of magnitude! In astrophysical environments, the density of photons is typically much larger than that of protons, unless there are residual masses from explosions (SNRs) or an accelerator with a molecular clouds and large rate of star formation (starbursts). Hence, even if the cross 10 section for pp interaction is about 2 orders of magnitude larger than that of pγ, this last may dominate.

11 Reminder: Reaction thresholds m p, p p m t,p t s=e cm 2 c=1 t + p M 1 + M M n f s = M f = E 2 2 tot p tot th The threshold of a reaction corresponds to the energy to produce all final states at rest. Remember also that from the invariance of total 4-momentum squared in the CM and Lab frame: total energies in the lab p s = q (E p + E t ) 2 ( ~p p + ~p t ) 2 = q m 2 p + m 2 t +2E p E t (1 p t cos ) At threshold and in the lab frame (the p rest-frame): p sth = q m 2 p + m 2 t +2E p m t = X f M f m 2 p + m 2 t + 2(E k,p + m p )m t =( X f M f ) 2 E k,p = (P f M f ) 2 (m p + m t ) 2 2m p

12 p-gamma Direct photo-production of pions: eg. p +! p + 0 p sth = m p + m = 1.08 GeV Energy of the photon in the lab: = m (m +2m p ) 2m p For delta-resonance Δ(1232) it is larger: m 2 p = m2 2m p 150MeV 340MeV If the proton is at rest (Ep = mp), ϵ= 340 MeV in the lab Hence in the CM: E p =γp mp and the photon energy is : 0 m 2 m 2 p = p = 2 m 2 p 2m p 2E 0 p m 2 p Hence the accelerated proton must have a threshold energy in the CM frame of: p-p E 0 p = 2 p m 2 m 2 p p 300GeV 1MeV 0

13 The end of the CR spectrum: GZK cut-off [Greisen 66; Zatsepin & Kuzmin66] 2.73 K 3k B T effective energy for Planck spectrum of CMB 0 =3k B T = ev = m2 2m p m 2 p 340MeV p = 0 = Hence the threshold energy of the proton in the CM is E p ~ γ p m p = ev Integrating over Planck spectrum E p,th ~ ev recent past now

14 Assumption on average energy fractions Kelner & Aharonian TeV 0.1 TeV

15 Two Body Decay Kinematics Each neutrino takes about 1/4 of the pion energy (on average) In the Lab (pion at rest)

16 Three Body Decay Kinematics In the LAB and for Eπ >> mπ (mass of electrons/neutrinos neglected)

17 p-gamma On average 1/3 of the p energy goes into pions p +! +! p + 0 2/3 p +! +! n + + 1/3 0! + ½ of pion energy goes into each photon +! µ + + µ! (e + e µ )+ µ On average 1/4 of the pion energy goes into each neutrino Assuming : dn p / Ep 2 => dn p / Ep 2 = E 2 de p de p Since : E E p = x 1 20 multiplicity and BR => de p = x 1 de x 2 and dn de / E 2 1 x dn / 2 1 de 3 E 2 dn p de p =2 2 3 E 2 x = x dn p de p =2 1 3 E 2 x = E E 2 => dn de = 1 4 dn de f

18 Multiplicities multiplicities from SOPHIA MC pp p-gamma Kelner & Aharonian

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20 Gamma-neutrino connection at source (before oscillations) dn de = dn de for p p dn de = 1 4 dn de for p What happens during propagation of messengers to us?

21 Oscillations in 3 families oscillation probability in 3 families (E. Lisi s lectures) solar U e1, U e2 θ 12 CHOOZ U e3 θ 13 atmospheric U e3 θ 13 U µ3,u τ3 θ 23 MNSP matrix

22 Astrophysical neutrino oscillations We can express the phase in astro units: For astrophysical source at kpc-distances emitting νs of 10 TeV: COSφ averages to zero since the extension of sources is about 1 pc and their distance is of the order of 1 kpc so the baseline is known with precision 1/1000 not 1/10 8 hence we use the incoherent term ν α \ν β ν e ν µ ν τ P(ν e ν e ) = 2 U ei Uei 2 = U i e1 4 + U e U e3 4 = = 0.56 ν e 60% 20% 20% ν µ 20% 40% 40% ν τ 20% 40% 40% P(ν e ν µ ) = 2 U ei Uµi 2 = U i e1 2 U µ1 2 + U e2 2 U µ U e 3 2 U µ1 2 = = 0.22 P(ν e ν τ ) = 2 U ei Uτi 2 = U i e1 2 U τ1 2 + U e2 2 U τ U e 3 2 U τ1 2 = = 0.22

23 From source to Earth At source: e : µ : 1:2:0 pion-muon decay ν α \ν β ν e ν e ν µ ν τ 60% 20% 20% ν µ 20% 40% 40% ν τ 20% 40% 40% At Earth: Maybe not true if charmed meson threshold is overcome: 0:1:0 (Sarcevic et al., arxiv: ) e : µ : 1:1:1 Maybe not true if muons cannot decay (Kashti & Waxman, PRL 95 (2005) ) 60% of νe survive and 2x20% come from 2xνµ = 100% 2x40%=80% of 2xνµ survive and 20% come from νe= 100% 20% of ντ come from νe and 2 x 40% from νµ= 100%

24 Neutrino events in a neutrino telescope

25 First flavor physics in a Neutrino telescope

26 Flavor ratio constrain the conditions at source e.g. magnetic fields (muon cooling), charm production, neutron dominated Future:

27 Neutrino rates (exercise)

28 Source Visibility IACT gamma-ray telescope South Pole is a special case Declination and zenith are complementary angles δ = 90 - θ [0, 180 ] latitude of the detector Warning! this plot applies to analyses using upgoing atmospheric neutrino dominated tracks but UHE neutrino astronomy has also to be done using down-going atmospheric muon dominated samples 28

29 The Sites and sky coverage North+South >60 o zenith 45 o -60 o 30 o -45 o North South

30 Concept of Neutrino Detector wavefront β 1 and θc 41 o θc θ 1.5deg E ν (TeV ) Between nm about 3.5 x 10 4 Cherenkov photons/m of a muon track Teresa Montaruli, Otranto, Sep 20-21,

31 Neutrino cross sections Energy transfer P. Mertsch s lecture l (k') Deep Inelastic Scattering (charged and neutral currents) with very large ν(k) momentum transfers. ~ Q 2 = - (k - k') (1 - y) E = E' = - (k - k') 2 ~M 2 W p y E ν 31

32 Large target mass due to small x-section µ + N! µ + X µ + N! µ + + X µ + N! µ + X µ + N! µ + X charged current neutral current At low energy, neutrino and anti-neutrino cross sections differ because the valence quark dominate hep-ph/

33 Muon/tau range and absorption/regeneration of neutrinos in the Earth P surv = e N A R dlcos (l) (E ) R = X 0 dx de de = de a + be = 1 a 0 E 0 0 E 0 de 1+ be /a = 1 b ln 1+ E 0 / E c ( ) Tau neutrinos never absorbed but loose energy L Abbate, TM, Sokalski, Astropart.Phys.23:57-63,2005 Bugaev, TM, Shlepin, Sokalski, Astropart.Phys. 21: ,2004

34 Effective Area 34

35 Neutrino selection & background rejection Upgoing thoroughgoing neutrino induced muons - Earth is a filter - or vertex identification of starting events (tracks and cascades)! μ Veto atmospheric muon tag " νμ atmospheric neutrino tag μ

36 Multi-messenger astrophysics

37 Reminder: Mean free path w = interaction prob. =w=n dx σ = cross section N= n. of target particles / volume P(x) = prob. that a particle does not interact after traveling a distance x P(x + dx) = prob that a particle has no interaction between x and x+dx = P(x+dx) = P(x) (1-wdx) P (x + dx) =P (x)+ dp dx = P (x) dx P (x)wdx dp P = wdx ) P (x) =P (0)e wx P(0) = 1 it is sure that initially the particle did not interacts = R xp (x)dx R P (x)dx = 1 w = 1 N Medium density I = = N c = Am p Atomic number in g/cm 2

38 Example: interaction length of CRs in the atmosphere I = = N c = Am p in g/cm 2 Total, inelastic and elastic (anti-)protons on proton cross section At p 10GeV/c σpp,inel 40 mb At 1 TeV σpp,inel 50 mb (1mb = cm 2 ) <Aatmosphere> = 14.5 (dominated by N) I = Am p pair For a beam of A nucleons: 2 1/ /3 2 σ nucl = π rn = π ( ro A ) 5 10 A cm = 50 mb x A 2/3 For p-air: 14.5*1.67 x kg /300 x cm 2 = 80 g/cm 2 For Fe-Air: 5 g/cm 2

39 The multi-messenger s horizons Proton horizon (GZK cut-off): p 2.7K! +! + n L = 1 p CMB n cm 2 400cm 3 10 Mpc The neutrino horizon is comparable to t observable universe! 1.95K! Z! X arxiv.org/pdf/ v2.pdf T. J. Weiler, Phys. Rev. Lett. 49, 234 (1982)

40 The proton horizon What about neutrons? for a neutron of E = 10 9 GeV = 10 6 TeV ldecay = Υcτ = E/mc 2 x c x 886 s = 10 9 GeV/1GeV x 3x10 8 m/s x 886s = 2.66 x m x 3.24 x kpc/m = 8.6 kpc

41 CR Composition

42 Deflection of CRs in B-field ) ( ) ( ) ( / / / / 2 G ZB ev E cm r ZeB pc r c ZevB r pv r mv = = = = 42 pc cm ev pc cm ev pc cm ev r ) ( ) ( ) ( = = = = = = =

43 Deflection of CRs in B-field Kelner, Aharonian, 43