Progresses in Neutrino Astronomy

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1 Progresses in Neutrino Astronomy Francesco Vissani INFN, Gran Sasso Theory Group Neutrino astronomy is a discipline aimed at investigating a class of high energy phenomena, by exploiting certain specific nuclear reactions. It enjoys a steadily improving instrumentation and observational successes, whereas theoretical investigations often suffer of an artificial separations among astrophysics, nuclear and particle physics. We review recent progresses in some selected fields emphasizing possible developments, making little attempts for completeness. With F. Aharonian, W. Fulgione, A. Ianni, G. Pagliaroli, F.L. Villante.

2 2/54. PRINCIPLES OF NEUTRINO TELESCOPES

3 3/54 Elastic scattering (ES) Figure 1: Due to the ES reaction, the (anti) neutrinos of several MeV hit the atomic electrons and produce events in the forward direction. We can distinguish ES events if we know the direction of the source. The sample of electrons from ES events locates a source of neutrinos, even when this is not identified by conventional astronomical means.

4 4/54 Figure 2: The Super-Kamiokande detector, located in the Kamioka mine (Japan). It reveals ultrarelativistic charged particles thanks to their Čerenkov-Vavilov radiation. Electrons and positrons above 5 MeV, produced by (anti) neutrinos, can be seen. Very high statistics, thanks to its 50 kton of H 2 O.

5 5/54 Figure 3: Image of the Sun obtained by exploing the electrons from the ES reaction as obtained by the Super-Kamiokande neutrino telescope. Similar results first obtained by Kamiokande and the detection of SN1987A are the reasons of 2002 Nobel prize awarded to Koshiba.

6/54 Inverse beta decay (IBD) Figure 4: IBD: an almost isotropic reaction occurring when E ν > 1.8MeV, with a much larger cross section then: σ IBD G 2 F E2 ν whereas σ ES G 2 F m ee ν.

6 6/54 Inverse beta decay (IBD) Figure 4: IBD: an almost isotropic reaction occurring when E ν > 1.8MeV, with a much larger cross section then: σ IBD G 2 F E2 ν whereas σ ES G 2 F m ee ν. This reaction is important, for free protons are abundant in H 2 O and C n H m targets. Positron is the footprint of the reaction; occasionally, even the neutron can be observed, providing a double tag of the event.

7 7/54 Exercise 1: By knowing 1. the Fermi constant G F = 10 5 GeV 2, 2. the useful convertion c = 200 MeV fm, 3. where the Fermi unit for nuclear size is fm=10 13 cm, estimate the cross section at 10 MeV (this is correct upto a factor of few). Find the free mean path in Earth s density, ρ 5 gr/cm 3.

occasionally hitted by neutrinos-or antineutrinos, even when their energy is below MeV.

8 8/54 Figure 5: The Borexino detector, located in the Gran Sasso laboratory (Italy). Its ultrapure scintillator permits us to observe the electrons-or positrons-which are occasionally hitted by neutrinos-or antineutrinos, even when their energy is below MeV. Borexino observes very well ν e : the neutron that follows IBD can be seen thanks to n + p D + γ(2.2 MeV).

9 9/54. Figure 6: Left, main panel: time distribution of SN1987A events, mostly due to IBD. Right, upper corner: Direction of the simulated events from a future galactic supernova: ES mostly in the center, IBD almost isotropical.

10 10/54 Neutral current (NC) deuterium dissociation Figure 7: By observing the neutron, it allows us to count all types of emitted neutrinos, thanks to the fact that Z-boson interacts in the same manner with all of them. The cross section is similar to IBD.

11 11/54 Figure 8: The SNO detector, formely located in Sudbury (now stopped). Using heavy water, and various ways to see the neutrons, it measured the total neutrino flux from the Sun.

12 . s -1 ) cm -2 6 (10! µ" SNO! ES SNO! CC SNO! NC! SSM ! e 6 (10-2 cm -1 s ) Figure 9: Electronic and non-electronic neutrino fluxes as measured by a variety of methods. NC s (blue oblique strip) measured by deuterium dissociation at SNO, tests and agrees with SSM predictions [SNO 2002]. 12/54

13 13/54. SOLAR NEUTRINOS

14 14/54. La catena PP (pp) 2p d + e + + ν e 2p + e d + ν e (pep) PPI 3 He + 3 He 4 He + 2p PPII d +p 3 He + γ 3 He + 4 He 7 Be + γ 3 He + p 4 He +e + + ν e PPIII (hep) (Be) 7 Be+e 7 Li + ν e 7 Be+p 8 B + γ 7 Li+p 2 4 He 8 B 2 4 He + e + + ν e (B) Figure 10: The main chain of nuclear energy generation in the Sun (Gamow 29). Directly tested by the relatively rare but energetic ν e from PPIII branch (Super-Kamiokande, SNO, Borexino) and by the more important PPII (Borexino). The PPI neutrinos are determined indirectly, i.e., from luminosity constraint.

15 15/54 Exercise 2: By knowing 1. the solar luminosity L = erg, 2. the energy of helium formation, Q = 27 MeV, 3. the solar distance, A.U.=150 million km, estimate the PP neutrino flux at Earth. Hint: simplify the PP chain to 4p α + 2ν e + Q, assume Q E ν.

16 16/54 Recent measurements of the PPIII neutrinos Figure 11: Boron-8 neutrinos flux. SNO uses ν e D e pp, SK=Super-Kamiokande and BX=Borexino use ES. Evidently, there is a disagreement with SSM predictions (which, let us repeat, agree with SNO measurements).

17 17/54 An energy dependent solar ν e deficit [PPIII, PPII and below] Figure 12: The solar neutrino conversion probability [Ianni 2010] The downturn at high energies is neatly explained by matter (MSW) effect, i.e., coherent scattering of electron neutrinos on solar electrons.

18 18/54 Different regimes of solar neutrino oscillations Assume a superposition of 2 mass states: ν e = cos θ ν 1 + sin θ ν 2 These 2 accumulate different and very large phases ϕ i 1; 8 < cos θ ν 1 e iϕ 1 + sin θ ν 2 e iϕ 2 below 1 MeV ν e, = : sin θ ν 2 e iϕ 2 above 10 MeV at high energy, an additional phase due to weak interactions with the medium implies ν e = ν 2, : this is the MSW effect. Thus 8 < P ee ν e ν e, 2 1 sin 2 2θ/2 0.6 Low Energy = : sin 2 θ 0.3 High Energy since θ = 34 ; the effect of the third neutrino is small.

19 19/54 Still a source of interesting theoretical problems! E.g.: There is an unexpected trend among the high energy neutrino data: they decrease toward the lowest energies. This is exciting (could be easily explained e.g., by a light sterile neutrino) though not very compelling. It emphasizes the need of more experimental information in the region between low and high energy solar neutrinos. Recent measurements of photospheric heavy element abundances lead to solar models unable to reproduce the helioseismic results. This solar composition problem demands a re-analysis of inputs and assumptions; impact of helioseismological constraints; opacity, distribution of the metals, role of secondary chain i.e., CNO neutrinos.

20 20/54 Figure 13: The counting spectrum in Borexino, with the various components of signal and background. Note the window for CNO detection, obscured by Bi β-decays. 210 Bi β-activity should be the same as 210 Po α-activity after a transient with time scale of 1 yr, since 210 Pb 32 yr 210 Bi 7 d 210 Po 200 d 206 Pb

21 21/54 ncno cpd 100ton n Β const. 2Σ 1Σ ncno cpd 100ton n Β const. 2Σ 1Σ Bi Bi Figure 14: Fit of 1 yr of data of Borexino, assuming an initial Po activity of 2000 cpd/100 ton and considering only the spectrum (left) or also the expected time evolution of the Polonium activity (right) [Villante 2011] It could be worthwhile to try already in Borexino, but surely in future detectors as SNO+ and LENA.

22 22/54. ANTINEUTRINOS from THE EARTH

23 23/54 ν e from terrestrial radioactivity 232 Th and 238 U beta-decay chains give rise to observable ν e [Eder 66; Marx 69]. The first type of nuclei are 4 times more abundant; positrons produced in the IBD reactions extend to 1.5 MeV and 2.5 MeV, respectively. ************** A hint of observation is due to KamLAND in 2005, 2008 in noisy conditions and concurrent ν e signal resulting from nuclear reactors. Past year, the first high confidence detection in a clean detector was obtained by Borexino, putting geo-neutrino science on firm observational bases.

24 24/54 Events/240p.e./252.6ton-year Borexino data best-fit reactors ν e contribution from geo-ν e background Light yield of prompt positron event [p.e.] Figure 15: The ν e signal is due in equal proportions to ν e s from far away reactors and those from terrestrial radioactivity: The latter component dominates at low energies. By future data, 232 Th and 238 U components should be distinguishable.

25 25/54. SUPERNOVA (anti)neutrinos

26 26/54 Formation of compact remnants Figure 16: Massive stars produce an inert iron core, which eventually collapses and produces a compact remnant, such as a neutrons star. The potential energy is radiated in neutrinos of all species.

27 27/54 Exercise 3: Estimate the thermal luminosity L R 2 T 4 by assuming R 10 km and T 5 MeV for all 6 (anti)neutrino types. Then, from the neutron star binding energy E G N M 2 /R (G N = cm 3 /(s 2 gr) and M = gr) find the time to form the neutron star. Finally calculate how many IBD events will take place in detector with mass 25 thousand tons of water if the supernova explodes at 10 kpc; assume E = 3T.

28 28/54 A thermal + initial accretion model compared with the sparse SN1987A data. Figure 17: Luminosity of electron ν e deduced from SN1987A; the peak is at 2.5 σ and resembles the expectations [Loredo 2002; Pagliaroli ].

29 29/54 The diffuse supernova neutrino background From the cosmic density rate of core collapse supernovae, R SN, and the average energy spectrum of ν e, dn/de, we calculate the diffuse flux: dn νe dt da de = c H 0 zmax 0 dz R SN (z) ΩΛ + Ω m (1 + z) dn 3 de (E(1 + z)) [Zel dovich & Guseinov 65; Ruderman 65] This is a steady signal: we do not need to wait our life to collect it! Note the cosmological redshift of the energy spectrum.

30 30/54 Figure 18: The 3 lines correspond to the uncertainty in the cosmic rate of SN explosion. The error-bars are the 2 σ ranges due to the uncertainties in the astrophysics of the explosion, as evaluated using SN1987A observations. Small rates! [FV 2011]

31 31/54 Appointment with the next galactic supernova vis MeV t sec Figure 19: Simulated events in Super-K for a supernova at 10 kpc [Pagliaroli 2009]. Using the times of the neutrino events (IBD+ES) and their directions (ES), one can infer with a precision of at least 10 ms the time of the gravity wave burst, comparable with its expected duration. A precious information for Virgo and LIGO!

32 32/54. HIGH ENERGY NEUTRINOS

33/54 Certain astronomical objects are thought to be point sources of very high energy ν, whose detection is one dream of the modern neutrino astronomy.

33 33/54 Certain astronomical objects are thought to be point sources of very high energy ν, whose detection is one dream of the modern neutrino astronomy. Figure 20: Principle of detection of (anti)ν µ : the induced muon flux. The Earth acts as µ-converter and as shield for atmospheric-µ. The prefential direction of observation is downward, where we have however atmospheric (anti)ν µ [Zheleitnik 57; Markov 60, Greisen 60].

34 34/54 Figure 21: The huge IceCUBE detector (South Pole). It reveals the (anti)µ occasionally produced by (anti)ν µ. It is optimized for high energy (>TeV) neutrinos.

Figure 23: The strict connection of γ and νµ produced in cosmic ray collisions can be tested, unless γ are

35 35/54 Figure 22: IceCUBE skymap, compatible with atmospheric νµ and Km3NET (Mediterranean) to be finalized in size and location. Figure 23: The strict connection of γ and νµ produced in cosmic ray collisions can be tested, unless γ are absorbed/modified. Standard astrophysical view: Observation of high-energy νµ is very difficult but perhaps possible and would amount to an unambiguous signal of cosmic ray collisions, and, hopefully, those in their source. F. Vissani Dublin, July 7, 2011

36 36/54 Potential neutrino sources and γ-rays Potential neutrino sources are characterized by their hadronic γ-rays distributed as I γ Eγ α e E γ /E c, with α = and E c = TeV PeV for π 0 and π ± are produced together. Figure 24: γ-ray intensities corresponding to a signal of 1 muon/km 2 yr above 1 TeV, evaluated assuming that the sources are transparent to their gamma rays [Aharonian 2011].

37 37/54 A few SNR s are known to satisfy the previous conditions!!! 11 Log Ν EΝ cm 2 s 1 TeV Log E Ν 1 TeV From the SNR RX J , we expect from HESS up to 2.4 signal events per year per km 2 and 1 background events for a detector with threshold set at 1 TeV [Villante 2008].

38 38/54 Exercise 4: Consider π 0 γγ in the laboratory frame. Show that the isotropic angular distribution of γ in the center of mass frame corresponds to a uniform distribution of E γ in the laboratory frame, i.e., dp π 0/dE γ = 2/E π 0 Derive the γ-flux dφ γ = 2 de γ E γ dφ π0 de π de π E π and show that it can be inverted obtain the π 0 -flux. Repeat the same steps for the ν µ -flux from π + µ + ν µ ; calculate it from the γ-flux, assuming π + π 0 as in pp-collision [FV06; Villante 08]

39 39/54. SUPERNOVA REMNANTS

40/54 Motivations for neutrinos from supernova remnants We want to test the hypothesis that the young SNR are the accelerators of the galactic cosmic rays and in some cases,

40 40/54 Motivations for neutrinos from supernova remnants We want to test the hypothesis that the young SNR are the accelerators of the galactic cosmic rays and in some cases, TeV neutrinos can help: Figure 25: The presence of a large number of targets yields mesons, e.g., pp π ± X, and thus to neutrinos, e.g., π + µ + ν µ [Drury, Volk, Aharonian 94].

41 41/54 SNR RX J : a case study Figure 26: Wang et al.: 393AD guest star=progenitor of RX J

42/54 Molecular clouds are present Figure 27: NANTEN (Sato et al., 2010) has observed the molecular clouds associated with RX J1713.

42 42/54 Molecular clouds are present Figure 27: NANTEN (Sato et al., 2010) has observed the molecular clouds associated with RX J studying the CO emissions, correlated with IR and anticorrelated with X ray emissions. The overdensities are named: peak A,B,C,D... Plausibly: overdensities formed by SN explosion and interacting with the shock wave. A,C,D most prominent. Peak C mass 400 M ; with 1 pc size ( 0.1 angular size) it has the column density of 20 µm of Lead.

43 43/54 TeV γ-ray emission is measured up to 100 TeV H.E.S.S. data 10.2 ) -1 TeV -2 s -1 dn/de ( cm Fit Fit 2004 Log EΓ 2.5 EΓ Energy ( TeV ) Log E Γ 1 TeV Figure 28: Thanks to HESS we know that the spectrum is non-trivial: it is well described by a broken-power-law or by a modified-exponential-cut. Power spectrum with 1.79 ± 0.06 and E c = 3.7 ± 1 TeV (Villante & FV 07)

44/54 Fermi view at GeV and above Figure 29: Counts > 3 GeV given by Fermi collaboration, claiming: a wide source in SNR location with spectrum E 1.5 γ from upper bound on γ of 0.

44 44/54 Fermi view at GeV and above Figure 29: Counts > 3 GeV given by Fermi collaboration, claiming: a wide source in SNR location with spectrum E 1.5 γ from upper bound on γ of GeV and measurements above; several point sources, including one sloping as E γ 2.45, outshining the wide source at GeV; diffuse background from the Milky Way. We superimposed the molecular clouds A, C, D of NANTEN.

45 45/54 An important result deserves comments and discussion [see Fukui 2011]: Leptonic (E γ e penetration (E γ p Eγ (γ+1)/2 ) or hadronic with energy dependent E γ+1/2 γ ), if primaries have γ 2 Lack of thermal X-rays: uniform medium+leptonic, or non-uniform. Seems important to measure the emission below 5 GeV. Eγ 1.7 not excluded firmly, that perhaps could still agree with usual hadronic emission and very efficient acceleration. Observational and theoretical progress expected in close future: Wait and see!

46 46/54 Any hope to consolidate the hadronic scenario? We should check the angular correlation with the Molecular Clouds, that interacting with CR s, give high energy γ. E (TeV) ON OFF σ Flux (cm 2 s 1 ) Range (TeV) (2 ± 0.6) (1 ± 0.4) (9 ± 3) (8 ± 2) (1 ± 0.5) (2 ± 2) ( ) HESS has 500 events above 30 TeV, but 70% is just noise. Probably we need to wait for a future γ-ray telescope measuring above 10 TeV needs, as CTA.

47 47/54. THE ULTIMATE FRONTIER

48 48/54 UHE energy neutrinos produce in certain media a coherent pulse of e leading to a GHz radio pulse: the Askaryan effect. END STATION A side view Approximately to scale crane hook: 13.7m 3m 13.4m 10.7m beamline 1.2m 2.4m 4.8m 10m 4m Figure 30: This effect in ice has been observed by the ANITA collaboration.

49/54 Figure 31: Moreover, at UHE also tau (not only muon or electron) neutrinos propagate for long distances and can lead to an observable signal. But which are the sources: AGN, GRB, or what?

49 49/54 Figure 31: Moreover, at UHE also tau (not only muon or electron) neutrinos propagate for long distances and can lead to an observable signal. But which are the sources: AGN, GRB, or what? The (even larger) uncertainty in the predictions would require more theoretical efforts. A guaranteed (weak) source of UHE neutrinos is due to the collisions of the UHECR with the microwave radiation field.

50 50/54 Figure 32: Present limits obtained by ANITA and RICE, based on the Askaryan effect, by Amanda (IceCUBE predecessor) and by the UHECR observatories. In the future, JEM-EUSO could contribute to the search (curves drawn by Medina-Tanco).

51 51/54. REMARKS and SUMMARY

52 52/54 Better speak of neutrino astronomies not astronomy Figure 33: Ranges (and optimal ranges) in log 10 (E/eV ) of the existing neutrino telescopes, spanning more than 10 decades in energy!

53 53/54 Neutrino astronomies seem to conform to usual evolution of astronomies: 1. Physics goal formulated; 2. Principle of detection (telescope); 3. First observation; 4. Systematic study; 5. Understanding. In this view, we note that the specific astronomies are in different stages: Solar neutrinos up to 4 (5?) pp; pep; CNO Geo-neutrinos up to 3 U / Th Supernova neutrinos up to 3 (4? 1?) first second HE neutrinos up to 2 (1?)??? Seems still a healthy field for theorists and experimentalists.

54 54/54 Neutrino astronomyies are stimulating and often considered branches of particle physics, but belong with the same / or more / rights to astronomy. Observations of low energy νs in very good shape, thanks to solar νs and to a new branch: geo-neutrinos. Theory of solar and supernova νs also proceeds, though slowly. Speculations and overoptimistic statements seem pretty common in / low and high energy / neutrino astronomy. High energy ν telescopes are a reality. Till now, no source seen yet but observational perspects are becoming more definite. There is space for progresses and perhaps for surprises. Many thanks for the attention!

Expectations for Galactic High Energy ν Sources

Expectations for Galactic High Energy ν Sources Francesco Vissani INFN, Gran Sasso Cosmic νs above TeV are the dream of modern neutrino astronomy. Here discussed: Expectations for galactic sources, with