Transcript SENSIBILITA’ DEL TELESCOPIO ANTARES PER NEUTRINI IN

Neutrino astronomy and telescopes
Crab nebula
Cen A
Teresa Montaruli, Assistant Professor, Chamberlin Hall,
room 5287, [email protected]
Overview
Neutrinos and their properties
Neutrino astronomy and connections to Cosmic rays and
gamma-astronomy
Neutrino sources and neutrino production
Neutrino telescopes
The Cherenkov technique and the photosensors
Search Methods
Current experimental scenario
Teresa Montaruli, 5 - 7 Apr. 2005
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 remedy……Unfortunately, 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 << a of
electromagnetic interactions
• 1956 Cowan and Reines : first detection of reactor neutrinos
by simultaneous detection of 2g‘s from e+ pair annihilation and
neutron
Teresa Montaruli, 5 - 7 Apr. 2005
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 -> e- + 37Ar, Eth=0.814 MeV, operated
continuously since 1970
Observed event rate of 2.56±0.23 SNU
(1 SNU =10-36 interactions per target atom per second)
Standard Solar Model prediction: 7.7+1.2-1.0 SNU 
Solar neutrino problem, now solved by oscillations
Teresa Montaruli, 5 - 7 Apr. 2005
Astrophysical neutrinos: from SN1987A
• SN1987A: 99% of binding energy in ns in a core collapse SN

Neutronization, ~10 ms 1051 erg e  p  n n e
Thermalization: ~10 s 31053 erg e  e n n
http://www.nu.to.infn.it/Supernova_Neutrinos/#7
ne  p  n  e
Teresa Montaruli, 5 - 7 Apr. 2005

The challenge
We learned:
Weak interactions make
neutrinos excellent probes
of the universe but their
detection is difficult !
Teresa Montaruli, 5 - 7 Apr. 2005
Neutrino Fluxes
Atmospheric ns
ns from WIMP
annhilation
Cosmic ns
Teresa Montaruli, 5 - 7 Apr. 2005
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 (particle=antiparticle) if only nL ad nR exist
The mass is a fundamental constant: needs to be measured!!

Direct neutrino mass measurements
ne< 3 eV ~ 3 x 10-9 mproton from b decay of 3H
(Z,A)(Z+1,A) + e- + ne
nµ< 0.17 MeV ~ 2 x 10-4 mproton from pmnm
nt< 18.2 MeV ~ 2 x 10-2 mproton from t 5pp0nt
Teresa Montaruli, 5 - 7 Apr. 2005
Neutrino mass =0?
Neutrino properties: oscillations
A n created in a leptonic decay of defined flavor is a linear superposition of mass
eigenstates
iE t
n a  Ua ,i n i
n
(
t
)

e
Ua ,i n i
2
2

a
E  p m
i
i
i
i
i
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
Pn a  n b    U a ,iU b* ,iU a* , jU b , j e
i, j
L=baseline
For 2 flavor:
 cos sin  

U  
  sin  cos 
 i mi2, j L
2E
where
Ea  p 
ma2
2p
2
2

1
.
27

m
(
eV
) L(km) 
2
2

Pn a  n b   sin 2 sin 
E (GeV )


Though oscillations are an indirect way of measuring the mass that requires
many different experiments to reach an understanding of the difference of the
Teresa
Montaruli, 5 - 7 they
Apr. 2005have the merit of being
square masses and of the flavors
involved,
sensitive to very small masses m2~<E>/L depending on the experiment design
The experimental scenario
Recent atmospheric neutrino experiments (Super-Kamiokande, MACRO,
Soudan 2) have demonstrated that the nm deficit is due to nm nt oscillations
with maximal mixing matm2 ~ 2.5 · 10-3 eV2 sin22q23 ~ 1
Solar neutrino experiments: (Cherenkov detectors: Super-Kamiokande, SNO)+
KAMLAND: scintillator detector looking for ne from reactors at ~180 km
average distance) deficit compatible with msun2 ~ 7.1 · 10-5 eV2 sin22q12 ~ 0.82,
could be due to ne nt or ne  nm
SNO at Sudbury Mine
3
2
1
Reactor neutrino experiments L~1 km (CHOOZ)
constrain the q13 mixing (no disappearance)
Teresa Montaruli, 5 - 7 Apr. 2005
Astronomy with particles



straight line propagation to point back to sources
Photons: reprocessed in sources and absorbed by extragalactic
backgrounds
For Eg > 500 TeV do not survive journey from Galactic Centre
Protons: directions scrambled by galactic and intergalactic magnetic fields
(deflections <1° for E>50 EeV)



d
Rgyro
 d  B 
1Mpc  1nG 
dB






E
E
0.1


20
 3 10 eV 
q
d
Rgyro
evB = mv2/Rgyro eB = p/Rgyro  1/Rgyro = B/E
Interaction length p + CMB 
+n
lgp = ( nCMB s ) -1 ~ 10 Mpc
Neutrons: decay gct  E/mn ct ~10kpc for E~EeV
g
p
Teresa Montaruli, 5 - 7 Apr. 2005
W49B
Messengers from the Universe
g
n
SN 0540-69.3
p
1 pc ~ 3 ly ~ 1018 cm
pgee
Crab
p + g p + n
3C279
Mrk421
Cas A
g+IRe+eg+radio
E0102-72.3
Local Group
g+MW
Gal Cen
<100 Mpc 1-109 TeV
Teresa Montaruli, 5 - 7 Apr. 2005
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
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
106 eV to ~1020 eV
~E-2.7
Balloons
satellites
knee
~E-3.1
EAS
Ankle:
1 km-2 century-1
~E-2.7
Teresa Montaruli, 5 - 7 Apr. 2005
After T. Gaisser, ICHEP02
CR acceleration at sources
Hillas Plot
The accelerator size
must be larger than
Rgyro
 B  R  18
10 eV

Emax  Z 
 1mG  1kpc 
energy losses in
sources neglected
R = acceleration site dimensions
Teresa Montaruli, 5 - 7 Apr. 2005
The knee
What is the origin of the knee?
1. Acceleration cutoff Emax~ZBL~Zx100TeV, change in acceleration process?
2. Confinement in the galactic magnetic field: rigidity dependent cut-off
3. 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
Teresa Montaruli, 5 - 7 Apr. 2005
Kascade - QGSJET
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 E~1019 eV could suggest a new light
population Protons are favored by all experiments.
AGASA: 111
scintillators + 27 m
detectors
Fe frac. (@90% CL): < 35% (1019 –1019.5 eV),
< 76% (E>1019.5eV)
Gamma-ray fraction upper limits(@90%CL)
34% (>1019eV) (g/p<0.45)
Teresa Montaruli, 5 - 7 Apr. 2005
56% (>1019.5eV) (g/p<1.27)
Relativity: 4-vectors
Controvariant pa = (E/c,px,pypz)=(E/c,p)
Covariant pa = (E/c,-px,-py-pz)
Also:
p  gm v
E  mgc 2
b  v/c
1
g
1 b 2
Scalar product of 2 4-vectors
qaka = (EqEk/c2-pqxpkx-pqypqy-pqzpqz)
Square
p2 = (E/c)2 - p2 = m2 = constant
Transform from one coordinate system to another
moving with speed v in the x direction (Lorentz
transformation)
p’x = g(px – bE/c)
p’y = py
p’z = pz
E’ = g(E – bpxc)
 E' / c  

 
 p' x  
 p'   
 y  
 p'  
 z  
g
bg
 bg
0
g
0
0
0
1
0
0
0
0  E / c 


0  p x 
0  p y 


1  p z 
In general: (E*,p*) in a frame moving at velocity bf:
Teresa Montaruli, 5 - 7 Apr. 2005
||=parallel to
direction of motion
T=transverse
Reaction Thresholds
mt,pp
mt,pt
t  p  M 1  M 2  ...  M f
2
2
s   M f  Etot  ptot
f
In the lab
Etot  mt  m p  Ek , p


ptot  p p
2
2mp Ek , p  p  Ek2, p
s = Ecm2
Energy of projectile
to produce particles
in the final state
at rest
c= 1

Etot  Ekin  m  p 2  m 2
2
2
2
Ekin  m  2m Ekin  p  m 2
True in any reference system
mp mt at rest
2

 M f

 f


  ( mt  m p  Ek , p ) 2  p 2p



 M f

 f


  ( mt  m p ) 2  Ek2, p  2(mt  m p ) Ek , p  p 2p



 M f

 f

  ( mt  m p ) 2  Ek2, p  2mt Ek , p


2
2
2


  M f   (mt  m p ) 2

Teresa Montaruli,
5 - 7
Apr. 2005
f


Ek , p 
2mt
Threshold for GZK cut-off
[Greisen 66;
Zatsepin & Kuzmin66]
Threshold pr p-gpN
2.73 K
Ek , p

 M f

f

g  p    pp 0
Eg  145 MeV
g  p    np
Eg  150 MeV


Energy of CMB photons:
2

  (mt  m p ) 2

(mN  mp ) 2  m 2p


2mt
2m p
in frame where p is at rest
=3kBT effective energy for Planck spectrum
And their energy in the proton rest frame is
Eg  g p g  150MeV
gp= 2· 1011 and the threshold
energy of the proton is then
Ep = gp mp = 2 ·1020 eV
Teresa Montaruli, 5 - 7 Apr. 2005
Integrating over Planck spectrum Ep,th~ 5 ·1019 eV
GZK cut-off?
[Greisen 66;
Zatsepin & Kuzmin66]
AGASA: 11 events, expects 2 @ E> 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: E>4 ·1019 eV




Air fluorescence detectors
HiRes 1 - 21 mirrors
HiRes 2 - 42 mirrors
Dugway (Utah)
Teresa Montaruli, 5 - 7 Apr. 2005
Anistropies
Galaxy cannot contain EHECR: at 1019 eV Larmour radius of CR p comparable to
Galaxy scale
AGASA: E>4 ·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 <1%) + 9 doublets (expect
1.7 chance probability <0.1%) at small scale (<2.5˚)
Not confirmed by HiRES
Triplet
close to
super-galactic
plane
See also
UHECR
correlation with
super-galactic
plane
astro-ph/9505093
Teresa Montaruli, 5 - 7 Apr. 2005
Neutrino production: bottom up
Beam-dump model: p0  g-astronomy p±  n-astronomy
Berezinsky et al, 1985
Gaisser, Stanev, 1985
p0
g
p
p
m nm
e n enm
n enmn t
Teresa Montaruli, 5 - 7 Apr. 2005
Neglecting g absorption
(uncertain) n  g
Targets: p or ambient g
From photon fluxes to n predictions:pp

Eg max
Eg min
Eg
dNg
dEg
dEg  K 
En max
En min
En
dNn
dEn
dEn
K = 1 pp
2 photons with
Eg 
Ep 
Ep
3
p  p  p  n p 
2

Ep
6
2nm and 1 ne with
En 
p  p  p  p p 0
Ep
Ep
4

Ep
12
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 ( =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
Teresa Montaruli, 5 - 7 Apr. 2005
Exercises!
From photon fluxes to n predictions: pg

Eg max
Eg min
Eg 
En 
Eg
dNg
dEg
dEg  K 
E p  x pp 
2
E p  x p p 
4
En max
En min
En
 10% E p
 5% E p
dNn
dEn
dEn
K = 4 pg
1) p  g    pp 0
2) p  g    np 
BR = 2/3
BR = 1/3
1) 2gs with 2/3× Eg = 2/3 ·0.1Ep
2) 2nms with 1/3× En = 1/3 0.1·Ep /2
 x pp  0.2
Teresa Montaruli, 5 - 7 Apr. 2005
K=4
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 v~c)
The probability of head-on encounters is
Magnetic
slightly greater than following collisions
Cloud
V
v
Head-on
Magnetic
inhomogenities
V
Following
This results in a net energy gain per
collision of

E
4 V 
  
E
3 c 
2
2nd order in the velocity of
the cloud
Teresa Montaruli, 5 - 7 Apr. 2005
1st order Fermi acceleration
2nd
The
order mechanism is a slow process.
The 1st order is more efficient since only
head-on collisions in shock waves
CasA Supernova Remnant in X-rays
Shock
fronts
High energy particles upstream and downstream
of the shock obtain a net energy gain when crossing
the shock front in a round trip
 
E
4 V 
  
E
3 c 
1nd order in the velocity of
the shock
Equation of continuity: r1v1=r2v2
For ionized gas r1/r2 = 4  v2 =4 v1
v1=u/4<v2
v2=u
John Hughe
upstream
Magnetic
Inhomogenities
v2>v1
Shock front at rest: upstream gas flows into shock
at v2 =u
And leaves the shock with v1 = u/4 Teresa Montaruli, 5 - 7 Apr. 2005
downstream
Blast shock
Interstellar medium
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
E
N
log P
log
log P
log b
1


E0
N0
N
E
log b
k


  
 dN  E
dE
log b
log P
N 0  E0 
v1 Particles lost in each round trip on the shock
P  1
log P
v1
1
c
3
3
 1
v v
log b
v1  v 2
1  v 2 / v1
b 2 1
c
log
Naturally predicts
dN
 E 2
dE
Teresa
5 - 7 Apr.
2005
CRs have steeper spectrum due
toMontaruli,
energy
dependence
of
diffusion in the Galaxy
v1 velocity of gas
leaving the shock
v2 velocity of gas
flowing into shock
dN
 E  2.7
dE