Astroparticle physics

Transcription

1 Astroparticle physics Big Bang and the particle physics and the dark size of the universe dark energy dark matter (axioins, sterile neutrinos, SUSY Q-balls, WIMPs, WIMPZILLAS, etc.) a smaller bang: the supernova (some of the above particles) Even smaller bangs: the cosmic rays (and some other particles) 1

2 Supernova: another lab of an astroparticle physicists Term supernova was coined by Baade and Zwicky in 1930s. Brighter than nova. Types of supernova Type I: absence of hydrogen lines Type Ia: no H line, strong Si line Type Ib: no H line, no Si line, strong He line Ic: no H line, no Si line, no (or weak) He line Type II: presence of hydrogen lines 2

3 Physics: Type Ia: probablythermonuclear explosion of a while dwarf accreting matter from a binary companion Type II: core collapse supernova 3

4 Chandrasekhar limit High-density stard can be stable only if their masses are less than M Ch. At high pressure matter degenerat Fermi gas. E F p 2 F /m e, for non-relativistic electrons E F p F, for relativistic electrons Consider an equilibrium configuration (Landau, 1932) N fermions in asphere of radius R. n N R 3 Volume per fermion: 1/n (Pauli principle) 4

5 By Heisenberg uncertainty principle, p F hn 1/3 n 1/3 E F { p 2 F n2/3 /m e N2/3 m e, for non relativistic electrons R 2 p F n 1/3 N1/3 R, for relativistic electrons Gravitational energy per fermion: E G GMm B R, where M = Nm B. NB: even though th epressure comes from electrons, all the mass is in baryons. 5

6 Total energy: E tot = E F + E G { N 2/3 m e GMm R 2 B R N 1/3 R GMm B R, for non relativ. for relativ. E tot = min equilibrium. In the non-relativistic case, ( 1 R R) 1 always has a minimum. 2 In the relativistic case, E tot = const R If const > 0, E tot can be decreased by increasing R until the system enters non-relativistic regime, the previous case. STABLE. If const > 0, E tot can be decreased without bound by decreasing R. COLLAPSE. 6

7 The critical mass is determined by setting to zero the coefficient of 1/R. E (N 1/3 GNm 2 B ) 1 R N 1/3 max GNm2 B = 0 ( 1 ) 3/2 = ( mpl ) 3 = ( ) 3 GeV = N max Gm 2 B m B 1GeV M max = N max m B = GeV M More precisely, M max = M Chandra = 1.4M 7

8 This depends only on fundamental constants! What is the radius? It is determines by the onset of degeneracy: E F N1/3 R where m is the mass of the fermion. = m R N1/3 max m Two special cases of interest: = 1 m ( ) GeV = GeV m normal matter: m = m e, R /m e 10 8 cm (white dwarf) nuclear matter m = m n, R /m n cm (neutron star) 8

9 A more precise calculation givess the following relation between the mass and the density: M stable (WD) unstable stable (NS) unstable ρ 9

10 Type Ia Supernova Thermonuclear explosion: WD made of CO burns into Fe. Total energy: E = ( m C 12 ) 56 m MWD Fe m C erg where m C, m Fe are the masses of 12 C and 56 Fe. Most energy is released in light, heat. Neutrinos carry only a negligible fraction of energy. 10

11 Core collapse supernova (Type II) A massive star, M > M burns H, He into Si core, into Fe core. When the mass of the Fe core exceeds the Chandrasekhar mass, the core collapses into a neutron star (ρ g/cm 3 ). Initial temperature MeV, so the weak interactions are important: pe nν e. Is there enough time to produce neutrinos? Compare τ weak,τ grav τ weak 1 σn σ G 2 F T 2, n ρ/m n 11

12 ( 10MeV τ weak s T ) ( 2 ρ ) g/cm 3 τ grav 1 GN ρ 0.4ms ( Since τ weak τ grav, produce a lot of neutrinos!! ρ g/cm 3 ) 1/2 12

13 Total energy: E total = GM2 R R=R NS R=R F e erg neutrinos erg kinetic energy of outgoing envelope erg light erg 13

14 Core collapse supernova A wonderful tool for setting bounds on particles with mass below 50 MeV (axions, sterile neutrinos). A place to discover new physics! SN1987A: neutrinos observed supernova asymmetries and the pulsar velocities. Discuss on the example of sterile neutrinos 14

15 Supernova asymmetries and new physics Weakly interacting particles, e.g. sterile neutrinos are emitted from a supernova with anisotropy Example: terile neutrinos with masses and mixing angles consistent with dark matter can explain the pulsar velocities [AK, Segrè; Fuller, AK, Mocioiu, Pascoli] 15

16 The pulsar velocities. Pulsars have large velocities, v km/s. [Cordes et al.; Hansen, Phinney; Kulkarni et al.; Lyne et al. ] A significant population with v > 700 km/s, about 15 % have v > 1000 km/s, up to 1600 km/s. [Arzoumanian et al.; Thorsett et al. ] 16

17 A very fast pulsar in Guitar Nebula HST, December 1994 HST, December

18 Map of pulsar velocities 18

19 Proposed explanations: asymmetric collapse [Shklovskii] (small kick) evolution of close binaries [Gott, Gunn, Ostriker] (not enough) acceleration by EM radiation [Harrison, Tademaru] (kick small, predicted polarization not observed) asymmetry in EW processes that produce neutrinos [Chugai; Dorofeev, Rodinov, Ternov] (asymmetry washed out) cumulative parity violation (it s not cumulative ) 19

20 Asymmetric collapse...the most extreme asymmetric collapses do not produce final neutron star velocities above 200km/s [Fryer 03] 20

21 Supernova neutrinos Nuclear reactions in stars lead to a formation of a heavy iron core. When it reaches M 1.4M, the pressure can no longer support gravity. collapse. Energy released: E G NMFe 2 core R 99% of this energy is emitted in neutrinos erg 21

22 Pulsar kicks from neutrino emission? Pulsar with v 500 km/s has momentum M v gcm/s SN energy released: erg in neutrinos. Thus, the total neutrino momentum is P ν; total g cm/s a 1% asymmetry in the distribution of neutrinos is sufficient to explain the pulsar kick velocities But what can cause the asymmetry?? 22

23 Magnetic field? Neutron stars have large magnetic fields. A typical pulsar has surface magnetic field B G. Recent discovery of soft gamma repeaters and their identification as magnetars some neutron stars have surface magnetic fields as high as G. magnetic fields inside can be G. Neutrino magnetic moments are negligible, but the scattering of neutrinos off polarized electrons and nucleons is affected by the magnetic field. 23

24 Core collapse supernova Onset of the collapse: t = 0 core Fe 24

25 Core collapse supernova Shock formation and neutronization burst : t = 1 10 ms PNS ν burst shock Protoneutron star formed. Neutrinos are trapped. The shock wave breaks up nuclei, and the initial neutrino come out (a few %). 25

26 Core collapse supernova Thermal cooling: t = s PNS ν thermal Most of the neutrinos emitted during the cooling stage. 26

27 Electroweak processes producing neutrinos (urca), p + e n + ν e and n + e + p + ν e have an asymmetry in the production cross section, depending on the spin orientation. The asymmetry: σ( e, ν) σ( e, ν) ǫ = g2 V g2 A g 2 + k k 0, V 3g2 A where k 0 is the fraction of electrons in the lowest Landau level. 27

28 In a strong magnetic field, B=3x10 16 G K B=10 G B=3x10 15 G µ, MeV k 0 is the fraction of electrons in the lowest Landau level. Pulsar kicks from the asymmetric production of neutrinos? [Chugai; Dorofeev, Rodionov, Ternov] 28

29 Can the weak interactions asymmetry cause an anisotropy in the flux of neutrinos due to a large magnetic field? No ν e ν e ν e ν e ν e νe ν e ν e ν e νe νe ν e ν e ν e ν e Neutrinos are trapped at high density. 29

30 Can the weak interactions asymmetry cause an anisotropy in the flux of neutrinos due to a large magnetic field? No Rescattering washes out the asymmetry In approximate thermal equilibrium the asymmetries in scattering amplitudes do not lead to an anisotropic emission. Only the outer regions, near neutrinospheres, contribute (a negligible amount). However, if a weaker-interacting sterile neutrino was produced in these processes, the asymmetry would, indeed, result in a pulsar kick! 30

31 Sterile neutrinos leave the star without scattering. Hence, they give the pulsar a kick. B ν s ν s ν e ν e ν e ν e ν e ν s ν s ν s Sterile neutrinos emitted asymmetrically by oscillations: [D Olivo, Nieves, Pal; Semikoz] 31

32 The bounds on sterile neutrinos m s (kev) 10 x ray limit pulsar kicks (allowed) excluded (x rays) for Ω =0.2 s Lyman α bound for production above 100 GeV dark matter produced via DW DW=Dodelson+Widrow; for discussion of the bounds, see hep-ph/ sin 2 θ 32

33 Other predictions of the pulsar kick mechanism Stronger supernova shock [Fryer, AK] Convection Region iii i ii 33

34 Other predictions of the pulsar kick mechanism Stronger supernova shock [Fryer, AK] 34

35 Other predictions of the pulsar kick mechanism Stronger supernova shock [Fryer, AK] No B v correlation is expected because the magnetic field inside a hot neutron star during the first ten seconds is very different from the surface magnetic field of a cold pulsar rotation washes out the x,y components Directional Ω v correlation is expected, because B the direction of rotation remains unchanged only the z-component survives 35

36 Gravity waves ν ν ν ν ν ν ν ν ν ν ν ν ν PULSAR Artist s conception by Roulet [Summer School lectures in Trieste] Rotating beam of neutrinos is the source of GW B 36

37 Gravity waves ν ν ν ν ν ν ν ν ν ν ν ν ν PULSAR Artist s conception by Roulet [Summer School lectures in Trieste] Rotating beam of neutrinos is the source of GW B Predicted correlation: direction of v and Ω. 37

38 Gravity waves at LIGO and LISA 1/2 θ, Hz noise signal at 1 kpc (f), Hz 1/2 TT θ Signal at 1 kpc Initial LIGO NB Adv LIGO WB Adv LIGO frequency, Hz [Loveridge, PR D 69, (2004)] frequency f, Hz 38

39 Physics of ultrahigh-energy cosmic rays Ultrahigh-energy cosmic rays Ultrahigh-energy neutrinos 39

40 Spectrum of cosmic rays 40

41 Spectrum of cosmic rays Probably galactic origin 41

42 Spectrum of cosmic rays Probably galactic origin Definite prediction 42

43 Spectrum of cosmic rays Probably galactic origin Greisen Zatsepin Kuzmin cutoff Definite prediction 43

44 Greisen-Zatsepin-Kuzmin cutoff Cosmic microwave background radiation has temperature 2.7K = 10 4 ev Protons interact with the CMBR and lose energy to pion photoproduction: Threshold: s = pγ nπ +, pπ 0 m 2 p + 2E γ E p > m p + m π or E p > ev Nucleon loses about 20% of its energy in each interaction Energy attenuation length: R GZK 50 Mpc 44

45 Cross section of pγ Nπ rises rapidly at the resonance 45

46 nucleon interaction length (dashed line) and energy attenuation length (solid line) 46

47 nucleon interaction length (dashed line) and energy attenuation length (solid line) Hubble distance 47

48 nucleon interaction length (dashed line) and energy attenuation length (solid line) Hubble distance GZK distance 48

49 expect a sharp drop in the flux of cosmic rays: E < ev should see sources from entire universe E > ev should see sources only within 50 Mpc 49

50 What about photons? λ Atten 100 Mpc 1 Mpc 10 kpc E γ (ev) 50

51 AGASA, HiRes experiments 1e+26 1e+25 J(E) x E 3 (m -2 sr -1 s -1 ev 2 ) 1e+24 1e+23 1e+22 1e+17 1e+18 1e+19 1e+20 1e+21 Energy (ev) AGASA events beyond the cutoff... 51

52 Air showers 52

53 Detection techniques Fluorescent detection: Fly s Eye, Hi Res, EUSO Surface array: AGASA Both: Pierre Auger

54 Pierre Auger 54

55 Pierre-Auger Observatory Northern Auger in Utah N#S Southern Auger in Argentina #S 1/22/ km Astronomy Colloquium at UCLA 39 55

56 56

57 Water tanks 57

58 Fluorescence detector 58

59 59

60 Lidar is used to monitor the atmosphere 60

61 Events (preliminary) PRELIMINARY 61

62 62

63 63

64 The GZK puzzle is two-fold what are the sources of UHECR? Why no GZK cutoff (assuming AGASA results are correct)? 64

65 Cosmic accelerators Can accelerate particles to ev? 65

66 Cosmic accelerators 66

67 Acceleration in AGN and radio galaxies 67

68 Fermi acceleration + B E 1 E 2 68

69 New physics? Supermassive relic particles Topological defects (cosmic strings, etc.) New exotic particles 69

70 Supermassive relic particles Could be copiously produced at the end of inflation, during reheating, from gravity or other interactions at the high scale. Can be the dark matter. Long lifetime is difficult to explain, but models exist [e.g., Kolb et al.; Kuzmin, Rubakov] Decays can produce UHECR in our galactic halo, hence no GZK cutoff 70

71 Superheavy dark-matter particles producing UHECR Observational signature: North-South asymmetry. 71

72 Topological defects High energy density in the core. Decays can produce high-energy particles. Long lifetimes. 72

73 AGASA vs HiRes 73

74 Pierre Auger will clarify the experimental situation Fitting even HiRes data alone might require exotic nearby sources 1e+26 1e+25 J(E) x E 3 (m -2 sr -1 s -1 ev 2 ) 1e+24 1e+23 1e+22 1e+17 1e+18 1e+19 1e+20 1e+21 Energy (ev) However, regardless of whether the GZK cutoff is there, there is great physics one can do with detectors like Pierre Auger and EUSO. 74

75 Extreme Universal Space Observatory (EUSO) EUSO is designed to observe fluorescent air showers initiated by extremely high energy cosmic rays - and neutrinos 75

76 EUSO is designed to observe fluorescent air showers initiated by extremely high energy cosmic rays - and neutrinos 76

77 Ultrahigh-energy neutrinos 77

78 UHE neutrinos are out there! Astrophysical sources pγ interactions imply a (predictable) flux of cosmogenic UHE neutrinos. Additional sources can provide additional flux. Hopefully, they will be discovered in the near future Pierre Auger, ANITA, ICE CUBE, EUSO, OWL,... 78

79 UHECR are deflected by magnetic fields + Virgo + Coma Centaurus + + Hydra + A Pavo + Perseus [Degrees] 79

80 UHE neutrino astronomy 80

81 Pierre Auger can detect UHE neutrinos ν

82 Certain predictions and limits 82

83 log 10 E 2 Φ ν (E) [cm 2 s -1 sr -1 GeV] Limits from AMANDA-B Charm D pred. Atmospheric neutrinos AMANDA-B10 E -2 ν e Frejus ν µ (diff. at E ν =2.6 TeV) MACRO E -2 ν µ Baikal E -2 ν e Charm D A-B10 lim. SS QC A-B10 lim. AMANDA-B10 E -2 ν µ pred log 10 (E ν ) [GeV] 83

84 Detection strategy relies on the knowledge of the neutrino-nucleon cross section at s 10 6 GeV Calculations of σ νn at ev necessarily use extrapolations of PDF and standard model parameters [Gandhi, Reno, Quigg, Sarcevic] Cross Section (cm 2 ) Standard Model, from Gandhi et al, 1998 Sigma_total, extra dimensions plus SM, M 4+n = 1.0 TeV Sigma_total, extra dimensions plus SM, M 4+n = 1.2 TeV Sigma_total, extra dimensions plus SM, M 4+n = 2.0 TeV Neutrino Energy (ev) 84

85 Several approved and proposed experiments plan to detect UHE neutrinos by observations of nearly horizontal air showers. Neutrinos are the only particles that interact weakly enough to produce horizontal air showers (assuming the cross section σ νn cm 2 at ev) Hence, particle ID is straightforward. How well do we know the neutrino-nucleon cross section? New physics contributions? Can saturation affect the cross section at high energy? Is it possible to measure the neutrino-nucleon cross section at these energies? 85

86 Calculations of neutrino-nucleon cross section at s 10 6 GeV Calculations of σ νn at ev necessarily use extrapolations of PDF and standard model parameters. Cross Section (cm 2 ) Standard Model, from Gandhi et al, 1998 Sigma_total, extra dimensions plus SM, M 4+n = 1.0 TeV Sigma_total, extra dimensions plus SM, M 4+n = 1.2 TeV Sigma_total, extra dimensions plus SM, M 4+n = 2.0 TeV Neutrino Energy (ev) 86

87 Calculations of neutrino-nucleon cross section at s 10 6 GeV Calculations of σ νn at ev necessarily use extrapolations of PDF and standard model parameters. Cross Section (cm 2 ) Standard Model, from Gandhi et al, 1998 Sigma_total, extra dimensions plus SM, M 4+n = 1.0 TeV Sigma_total, extra dimensions plus SM, M 4+n = 1.2 TeV Sigma_total, extra dimensions plus SM, M 4+n = 2.0 TeV Neutrino Energy (ev) 87

88 Calculations of neutrino-nucleon cross section at s 10 6 GeV Calculations of σ νn at ev necessarily use extrapolations of PDF and standard model parameters. Cross Section (cm 2 ) Standard Model, from Gandhi et al, 1998 Sigma_total, extra dimensions plus SM, M 4+n = 1.0 TeV Sigma_total, extra dimensions plus SM, M 4+n = 1.2 TeV Sigma_total, extra dimensions plus SM, M 4+n = 2.0 TeV Neutrino Energy (ev) 88

89 Neutrino-nucleon cross section at s 10 6 GeV Calculations necessarily use extrapolations of PDF and standard model parameters. Cross Section (cm 2 ) Standard Model, from Gandhi et al, 1998 Sigma_total, extra dimensions plus SM, M 4+n = 1.0 TeV Sigma_total, extra dimensions plus SM, M 4+n = 1.2 TeV Sigma_total, extra dimensions plus SM, M 4+n = 2.0 TeV Neutrino Energy (ev) SM calculation [Quigg et al.] is most likely right, but we want to measure this cross section. 89

90 If the cross section is smaller, the Earth becomes more transparent to neutrinos. More neutrinos can get through the Earth, interact just below the surface and produce a charged lepton that originates an up-going air shower (UAS). The increase in UAS rate compensates for the decrease in HAS. The comparison of the two rates allows a measurement of the cross section at GeV Angular distribution of UAS can provide an additional independent information about the cross section 90

91 τ ν The probability of a neutrino conversion into an up-going τ grows with the mean free path λ ν, for λ ν < R, because the shadowing by the Earth decreases. 91

92 ν τ UAS requires a neutrino to interact and produce a τ below the surface. 92

93 ν τ UAS requires a neutrino to interact and produce a τ below the surface. The number of UAS is higher for a smaller cross section. 93

94 The shower probability per incident neutrino: UAS HAS σ, cm The energy threshold for detection of UAS was assumed E th = ev for curve 1 and E th = ev for curve 2. Additional UAS events, not included here, can be detected by EUSO or OWL via Cerenkov radiation of tau leptons. 94

95 In addition, the angular distributing depends on the cross section. 90 o o o σ= 1.0 x 10 σ= x o o 6.24 o Most probable UAS corresponds to chord length close to mean free path σ= 6.0 x σ= 3.90 x σ= 1.75 x σ= 1.61 x

96 Can one disentangle the unknown flux from the unknown cross section? Rate of HAS: F ν σ νn Rate of UAS: F ν /σ νn Angular distribution: cos θ peak 1/σ νn can determine σ νn and F ν independently All one needs is enough statistics. Need neutrino telescopes. Possible to do particle physics experiments using cosmic ray detectors 96

97 Conclusion Astrophysical observations can yield information about particle physics. These data are complementary to the collider experiments. Knowledge of astrophysical objects is important. in contrast with traditional astrophysics, particle physicists ask different questions about the same objects. Dark matter implies that new physics is out there to be discovered. Other signs of new physics (baryogenesis, cosmic rays, etc.). 97