Neutrino astronomy and telescopes

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Neutrino astronomy and telescopes
Crab nebula
Cen A
Teresa Montaruli, Assistant Professor, Chamberlin
Hall, room 5287, tmontaruli_at_icecube.wisc.edu
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Overview
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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
  remedyUnfortunately, 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-decay theory

Week interactions GF ltlt a of electromagnetic
interactions

• 1956 Cowan and Reines first detection of
  reactor neutrinos
• by simultaneous detection of 2gs from e pair
  annihilation and
• neutron

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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 -gt
e- 37Ar, Eth0.814 MeV, operated continuously
since 1970 Observed event rate of 2.560.23 SNU
(1 SNU 10-36 interactions per target atom per
second)  Standard Solar Model prediction
7.71.2-1.0 SNU ? Solar neutrino problem, now
solved by oscillations   
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Astrophysical neutrinos from SN1987A
http//www.nu.to.infn.it/Supernova_Neutrinos/7
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The challenge
We learned Weak interactions make neutrinos
excellent probes of the universe but their
detection is difficult !
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Neutrino Fluxes
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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 (particleantiparticle) if only nL ad
  nR exist
• The mass is a fundamental constant needs to be
  measured!!

Direct neutrino mass measurements nelt 3 eV 3 x
10-9 mproton from b decay of 3H (Z,A)?(Z1,A)
e- ne nµlt 0.17 MeV 2 x 10-4 mproton from
p?mnm ntlt 18.2 MeV 2 x 10-2 mproton from t
?5pp0nt Neutrino mass 0?
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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
Lbaseline 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 Dm2ltEgt/L
depending on the experiment design
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The experimental scenario
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Astronomy with particles

• straight line propagation to point back to
  sources
• Photons reprocessed in sources and absorbed by
  extragalactic backgrounds
• For Eg gt 500 TeV do not survive journey from
  Galactic Centre
• Protons directions scrambled by galactic and
  intergalactic magnetic fields (deflections lt1
  for Egt50 EeV)
• Interaction length p gCMB ? p n
• lgp ( nCMB s ) -1 10 Mpc
• Neutrons decay gct ? E/mn ct 10kpc for EEeV

q
d
Rgyro
evB mv2/Rgyro? eB p/Rgyro ? 1/Rgyro B/E
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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
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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
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CR acceleration at sources
The accelerator size must be larger than Rgyro
energy losses in sources neglected
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The knee

• What is the origin of the knee?
• Acceleration cutoff EmaxZBLZx100TeV, 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
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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 E1019 eV could suggest a new
light population Protons are favored by all
experiments.
Fe frac. (_at_90 CL) lt 35 (1019 1019.5 eV), lt
76 (Egt1019.5eV) Gamma-ray fraction upper
limits(_at_90CL) 34 (gt1019eV) (g/plt0.45)
56 (gt1019.5eV) (g/plt1.27)
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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) px g(px bE/c) py
py pz pz E g(E bpxc)
Also
In general (E,p) in a frame moving at velocity
bf
parallel to direction of motion Ttransverse
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Reaction Thresholds
Energy of projectile to produce particles in the
final state at rest
mt,pt
mt,pp
s Ecm2 c 1
True in any reference system
In the lab
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Threshold for GZK cut-off
Greisen 66 Zatsepin Kuzmin66
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

Integrating over Planck spectrum Ep,th 5 1019 eV
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GZK cut-off?
Greisen 66 Zatsepin Kuzmin66
AGASA 11 events, expects 2 _at_ Egt 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 Egt4 1019 eV

• Air fluorescence detectors
• HiRes 1 - 21 mirrors
• HiRes 2 - 42 mirrors
• Dugway (Utah)

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Anistropies
Galaxy cannot contain EHECR at 1019 eV Larmour
radius of CR p comparable to Galaxy scale AGASA
Egt4 1019 eV no evidence of anisotropies due to
galactic disc but large scale isotropy ? EHECR
are extra-galactic AGASA 67 events cluster ? 1
triplet (chance prob lt1) 9 doublets (expect
1.7 chance probability lt0.1) at small scale
(lt2.5) Not confirmed by HiRES
Triplet close to super-galactic plane See
also UHECR correlation with super-galactic
plane astro-ph/9505093
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Neutrino production bottom up
Beam-dump model p0 ? g-astronomy p ?
n-astronomy
Neglecting g absorption (uncertain) ?n ? ?g

Targets p or ambient g
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From photon fluxes to n predictionspp
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
Exercises!
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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 0.1Ep /2
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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
vc) 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
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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 r1v1r2v2 For ionized gas
r1/r2 4 ?? v2 4 v1
upstream
v1u/4ltv2
v2u
Shock front at rest upstream gas flows into
shock at v2 u And leaves the shock with v1 u/4
downstream
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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