Transcript Ice Fishing for Neutrinos - International Centre for

Ice-fishing for Cosmic
Neutrinos
Subhendu Rakshit
TIFR, Mumbai
Goals of neutrino astronomy
• Astrophysics:
To explore astrophysical objects like AGN or GRBs. Find
out sources of high energy cosmic rays. Main aim..
• Particle physics:
To explore beyond standard model physics options
which may affect neutrino nucleon cross-sections at high
energy. Other possibilities… Appeared in US particle
physics roadmap!
First step: To determine the incoming neutrino flux
Astrophysical motivations
• Historically looking at the same astrophysical
object at different wavelengths revealed many
details regarding their internal mechanisms
• A 3-pronged approach involving conventional
photon astronomy, cosmic ray astronomy and
neutrino astronomy will yield better results
Conventional astronomy with photons
• Ranges from 104 cm radio-waves to 10-14 cm high
energy gamma rays
• Pros:
 Photons are neutral particles. So they can point back to their
sources
 photons are easy to detect as they interact
electromagnetically with charged particles
• Cons:
 Due to the same reason they get absorbed by dust or get
obstructed
 Very high energy photons on its way interact with cosmic
microwave background radiation and cannot reach us
Cosmic ray astronomy
• Very high energy cosmic rays (protons, heavy nuclei,..)
do reach us from the sky
• It is difficult to produce such energetic particles in the
laboratory
• It is puzzling where they are produced and how they get
accelerated to such energies!!
• Although they can be detected on Earth, it is not possible
to identify the sources as their paths get scrambled in
magnetic fields  A serious disadvantage!
• Only very high energy(>1010 GeV) cosmic rays point
back to their sources
Neutrino astronomy
• The suspected sources of very high energy
photons and cosmic rays are believed to be the
sources of neutrinos as well
• Pros: Neutrinos being weakly interacting reaches
Earth rather easily
• Cons: Due to the same reason it also interacts
rarely with the detector material ⇒ Large
detector size!!
• Successful neutrino astronomy with the sun and
supernova. Now it is time to explore objects like
Active Galactic Nuclei or Gamma Ray Bursts
• Impressive range for future neutrino telescopes:
102 GeV to 1012 GeV!
Neutrino detectors
Underground
Air shower
Underwater / ice
GeV
TeV
1 PeV = 106 GeV
1 EeV = 109 GeV
PeV
EeV
Why a Km3 detector?
• Estimations of the expected amount of UHE neutrinos
can be made from the observed flux of cosmic rays at
high energies. This limits the size of the detector
• However such estimations are quite difficult as many
assumptions go in
• There can be hidden sources of neutrinos!!
• So the neutrino flux can always be higher!
o1
K
M
^
3
IceCube
• A Km3 detector
• PMTs detect Cherenkov
light emitted by charged
particles created by
neutrino interactions

νμ
The Cherenkov cone needs to be reconstructed to determine the
energy and direction of the muon
Used for calibration,
background rejection and airshower physics
- The predecessor of IceCube
IceCube is optimised for detection of muon
neutrinos above 1 TeV as:
• We get better signal to noise ratio
• Neutrino cross-section and muon range increases with
energy. Larger the muon range, the larger is the effective
detection volume
• The mean angle between muon and neutrino decreases
with energy like 1/√E, with a pointing accuracy of about 1◦
at 1 TeV
• The energy loss of muons increases with energy. For
energies above 1 TeV, this allows us to estimate the muon
energy from the larger light emission along the track
Detection strategy

• Cosmic rays produce muons in our
atmosphere, which can fake a
neutrino-induced muon signal
 background
• So we use the Earth to filter them
out!
• Upto PeV neutrinos can cross the
Earth to reach IceCube
• For high energy neutrinos Earth
becomes opaque as the probability
that the neutrinos will interact
becomes higher with  energy
IceCube
• So very high energy neutrinos can
reach Icecube only from the sky or
from horizontal directions!
Sources of neutrinos
• Signal: The neutrinos from astrophysical sources: AGN
or GRBs for example
• Background: Atmospheric neutrinos. They are
produced from cosmic ray interactions with the
atmosphere  A guaranteed flux well measured in
AMANDA. Agrees with expectations.
As the ATM  flux falls rather rapidly(∝ E-3) with energy,
at higher energy we can observe the ‘signal’ neutrinos
from AGN or GRBs free of these background neutrinos
Neutrino spectra
Note: At higher energies
the flux is smaller. But
higher energy neutrinos
also have higher crosssection. So detection
probability is also higher!
Another background
• Cosmogenic or GZK neutrinos:
UHE cosmic ray protons interact with CMBR photons to
produce these neutrinos via charged pion decay
However at IceCube the rate would be quite small
Eliminating backgrounds
• Energy cuts
• Directional cuts
• Directional signals
• Temporal considerations
• Production at astrophysical sources:
Initial flavour ratio
νe :νμ :ν τ =1:2:0
• Propagation through space:
Massive neutrinos undergo quantum mechanical oscillations.
So neutrinos reach Earth with a flavour
ratio ν :ν :ν =1:1:1
e
μ
τ
• Propagation through the Earth:
Neutrinos while propagating may interact with the Earth. CC or
NC interactions. τ propagation is more elaborate: τ→τ→
τ→τ...
• Detection at IceCube:
Muon neutrinos produce muons via CC interactions. All
neutrinos produce showers through NC interactions. A CC
interaction by a τ may produce spectacular signatures!
Production at astrophysical sources:
A proton gets accelerated and hits another proton or a
photon. They produce neutron, π+ and π0.Their
decay produces cosmic rays, neutrinos and photons
respectively
p +  → π+ + n
μ + νμ
+
e+ + ν e + ν μ
p +  → π0 + p
γ+γ
νe :νμ :ν τ =1:2:0
Propagation through space:
• For massive neutrinos flavour and mass eigenstates are
different. This implies that a neutrino of a given flavour
can change its flavour after propagating for sometime!
For example: µ ↔ e
Neutrino oscillation
At time t=0, we produce a e
νe (0) = a ν1 + b ν 2
After sometime t, the mass eigenstates evolve differently
νe (t) = a e-iE1t ν1 + b e-iE2t ν2
So the probability of detecting another flavour is nonzero
• Now remember the initial flavour ratio at source was
νe :νμ :ν τ =1:2:0
At source
• Recent neutrino experiments have established that
neutrino flavour states µ and τ mix maximally
• Hence it is of no wonder that after traversing a long
distance these two states will arrive at equal proportions
νe :νμ :ν τ =1:1:1
On Earth
• Note that although there were no tau neutrinos at the
source, we receive them on Earth!
Propagation through the Earth:
• While traversing through the Earth, neutrinos can
undergo
 a charged current(CC) interaction with matter. The neutrino
disappears producing e or mu or tau. The dominant effect
 or a neutral current interaction(NC) with matter. The neutrino
produces another neutrino of same flavour with lower energy
• As a consequence, the number of neutrinos decrease as
they propagate through the Earth.
• This depends on the energy of the neutrino. Higher
energy neutrinos get absorbed more, their mean free
path is smaller
int
int 
1
N
N A tot
µ detection
• Muons range: few Kms at TeV and tens of Km at EeV
• The geometry of the lightpool surrounding the muon
track is a Km-long cone with gradually decreasing radius
• Initial size of the cone for a 100TeV muon is 130m. At the
end of its range it reduces to 10m.
• The kinematic angle of µ wrt the neutrino is µ is
1◦/√(E/1TeV) and the reconstruction error on the muon
direction is on the order of 1◦
• Better energy determination for contained events. More
contained events at lower energy
~ Km long muon tracks
from µ
~ 10m long cascades
from e, τ
e detection
• In a CC interaction, a e deposits 0.5-0.8% of their energy
in an EM shower initiated by the electron. Then a shower
initiated by the fragments of the target
• The Cherenkov light generated by shower particles spreads
over a vol of radius 130m at 10TeV and 460m at 10EeV.
Radius grows by ~50m per decade in energy
• Energy measurement is good. The shower energy
underestimates the neutrino energy by a factor ~3 at 1 TeV
to ~4 at 1 EeV
• Angle determination poor! Elongated in the direction of e
so that the direction can be reconstructed but precise to
~10◦
τ detection
• The propagation mechanism of a tau neutrino is
different, as tau may decay during propagation
τ
τ
τ
τ
• As a result the tau neutrino never disappears.
For each incoming τ another τ of lower
energy reaches the detector
• The Earth effectively remains transparent even
for high energy tau neutrinos
• Tau decays produce secondary flux of e and µ
• Double bang events: CC interaction of τfollowed by
tau decay
• Lollipop events: second of the two double bang
showers with reconstructed tau track
• Inverted lollipop events: first of the two double bang
showers with reconstructed tau track. Often confused with a
hadronic event in which a ~100GeV muon is produced!
• For Eτ< 106 GeV, in double bang events showers are
indistinguishable. For Eτ~ 106 GeV, tau range is a few
hundred meters and the showers can be separated.
For 107 GeV < Eτ< 107.5 GeV, the tau decay length is
comparable to the instrumented detector vol. lollipop
Eτ> 107.5 GeV tau tracks can be confusing
Propagation equation of µ
N
d

d ( E , X )
1
dy
NC ( E y , y )

 ( E, X )  N A 
 ( Ey , X )
dX
int ( E )
1 y
dy
0
1
int 
1
N
N A tot
E
Ey 
1 y
Propagation equations of τ
 ( E, X )
X
1
 ( E, X ) 1 dy NC
dy


K  ( E y , X )  
K ( E, y) ( E y , X )
 ( E ) 0 1  y
1 y
0
 ( E , X )
 ( E, X )
1
dy CC
 

K ( E , y ) ( E y , X )

ˆ
X
 ( E ) 0 1  y
 ( E )
1
 
1
N
tot
A  N

dec
1
1
1
 CC  dec
ˆ

 
 E
( E, X , )  
 m

 c  

E
Ey 
1 y
 
CC
1
N A CC
N
NC ,CC
d

( Ey , y)
1
N
NC ,CC
K
( E , y )  tot
 N ( E )
dy
KCC ( E, y) 
dec
K
1
totN ( E )
dCC
N ( Ey , y)
dy
1 d   X ( Ey , y)
( E, y)  tot
 ( E )
dy
K ( E, y) 
1
1
KCC ( E, y)  dec
Kdec ( E, y )
 ( E )
 ( E )
Without energy
loss
Including
energy loss
Rakshit, Reya, PRD74,103006(2006)
Characteristic bump
Expected muon event rate per year at IceCube
µ induced
µ+ τ induced
Imprinted Earth’s
matter profile
• Production at astrophysical sources:
Initial flavour ratio
νe :νμ :ν τ =1:2:0
?
• Propagation through space:
Massive neutrinos undergo quantum mechanical oscillations.
So neutrinos reach Earth with a flavour
ratio
νe :νμ :ν τ =1:1:1
??
• Propagation through the Earth:
Neutrinos while propagating may interact with the Earth. CC or
NC interactions. τ propagation is more elaborate: τ→τ→
τ→τ...
• Detection at IceCube:
N xsection sensitive
Muon neutrinos produce muons via CC interactions. All
neutrinos produce showers through NC interactions. A CC
interaction by a τ may produce spectacular signatures!
• Detection of atm µs will enable us to probe CPTV,
LIV,VEP which change the standard 1/E energy
dependence of osc length. Due to high threshold of
IceCube, osc of these high energy atm neutrinos is less
• N xsection can get enhanced in XtraDim models
• N xsection can get reduced at high energies in color
glass condensate models
• Visible changes in muon rates, shower rates
• For xtradim upgoing neutrinos get absorbed at some
energy and also downgoing for higher energies
• For lower N xsection models angular dependence and
energy dependence for upgoing events are more important
• Crude neutrino flux determination from up/down events
• OK for fixed power flux, but otherwise contained muon
events are better. But poorer statistics
• Auger is better for UHE neutrinos. New physics effects
will be more dramatic
• IceCube can probe neutrino spectrum better as Xsection
uncertainties are only at high energies where the flux is
smaller
• Flavour ratio determination possible at IceCube as
different flavours have distinctive signatures.
Other possibilities
•
•
•
•
•
•
DM detection: Neutrinos from solar core
SUSY search: look for charged sleptons
RPV, Leptoquarks
Part of supernova early detection system!
New physics interactions at the detector
New physics during propagation
Summary
• UHE neutrinos: particle physics opportunities for the
future
• IceCube is a discovery expt.
• Determining neutrino spectrum independent of new
physics poses a challenge
• Even crude measurements at IceCube may provide
some clue about drastically different new physics
scenarios at high energies
• Some success with IceCube will lead to bigger detectors
• At present we just need to detect an UHE neutrino event
at IceCube!
Particle physics motivations
LHC CM energy ECM = 14 TeV
ECM
⇒
 E
 2M N E  14  17
 10 eV
LHC: E=108 GeV

TeV

Tevatron: E=106 GeV
Here we talk about neutrino flux of 1012 GeV!
⇒ ECM = 14 ×100 TeV
N cross-sections
• We need PDF’s for x < 10-5 for E>108
GeV
MW2
103
x
2M N E E / GeV
• Several options but not much discrepancy!
• GRV and CTEQ cross-sections differ at
the most by 20%
Beacom et al, PRD 68,093005(2003)
e shower(CC+NC)
For downgoing μ
Horizontal μcreating a
detectable μ track
τlollipop
τdouble
bang