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Neutrino astronomy and telescopes

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### Neutrino astronomy and telescopes

Crab nebula

Cen A

Teresa Montaruli, Assistant Professor, Chamberlin Hall, room 5287, [email protected]

Neutrinos and their properties

Neutrino astronomy and connections to Cosmic rays and gamma-astronomy

Neutrino sources and neutrino production

Neutrino telescopes

Search Methods

The Cherenkov technique and the photosensors

Current experimental scenario

OverviewTeresa 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-decaytheory

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

- SN1987A:99% of binding energy in ns in a core collapse SN Neutronization, ~10 ms 1051 erg
- Thermalization: ~10 s 31053 erg

http://www.nu.to.infn.it/Supernova_Neutrinos/#7

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

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 pmnm

nt< 18.2 MeV ~ 2 x 10-2 mproton from t5pp0nt

Neutrino mass =0?

Teresa Montaruli, 5 - 7 Apr. 2005

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

L=baseline

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 Dm2~<E>/L depending on the experiment design

Teresa Montaruli, 5 - 7 Apr. 2005

Recent atmospheric neutrino experiments (Super-Kamiokande, MACRO,

Soudan 2) have demonstrated that the nm deficit is due to nm nt oscillations with maximal mixing Dmatm2~ 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 Dmsun2~ 7.1 · 10-5 eV2 sin22q12 ~ 0.82, could be due to nent or ne nm

Reactor neutrino experiments L~1 km (CHOOZ) constrain the q13 mixing (no disappearance)

3

2

1

SNO at Sudbury Mine

The experimental scenarioTeresa Montaruli, 5 - 7 Apr. 2005

straight line propagation to point back to sources MACRO,

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)

Interaction length p + gCMBp + n

lgp =( nCMBs ) -1~10 Mpc

Neutrons: decay gct E/mn ct ~10kpc for E~EeV

Astronomy with particlesq

d

Rgyro

evB = mv2/Rgyro eB = p/Rgyro 1/Rgyro = B/E

Teresa Montaruli, 5 - 7 Apr. 2005

W49B MACRO,

g

n

SN 0540-69.3

p

p + gp + n

pgee

Crab

3C279

Mrk421

g+IRe+e-

Cas A

Local Group

g+radio

E0102-72.3

g+MW

Gal Cen

<100 Mpc 1-109 TeV

Messengers from the Universe1 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

Teresa Montaruli, 5 - 7 Apr. 2005

10 MACRO, 6 eV to ~1020 eV

~E-2.7

~E-3.1

knee

Balloons

satellites

EAS

Ankle:

1 km-2 century-1

~E-2.7

After T. Gaisser, ICHEP02

The CR spectrumSN 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

Teresa Montaruli, 5 - 7 Apr. 2005

Hillas Plot MACRO,

R = acceleration site dimensions

CR acceleration at sourcesThe accelerator size

must be larger than

Rgyro

energy losses in sources neglected

Teresa Montaruli, 5 - 7 Apr. 2005

Kascade - QGSJET MACRO,

The knee- What is the origin of the knee?
- Acceleration cutoff Emax~ZBL~Zx100TeV, 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

Teresa Montaruli, 5 - 7 Apr. 2005

AGASA: 111 MACRO, scintillators + 27 m detectors

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.

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) 56% (>1019.5eV)(g/p<1.27)

Teresa Montaruli, 5 - 7 Apr. 2005

Controvariant p MACRO, a = (E/c,px,pypz)=(E/c,p)

Scalar product of 2 4-vectors

qaka = (EqEk/c2-pqxpkx-pqypqy-pqzpqz)

Square

p2 = (E/c)2 - p2 = m2 = constant

Relativity: 4-vectorsCovariant pa = (E/c,-px,-py-pz)

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)

Also:

In general: (E*,p*) in a frame moving at velocity bf:

||=parallel to direction of motion

T=transverse

Teresa Montaruli, 5 - 7 Apr. 2005

Energy of projectile MACRO,

to produce particles

in the final state

at rest

mp

mt at rest

Reaction Thresholdsmt,pt

mt,pp

s = Ecm2 c= 1

True in any reference system

In the lab

Teresa Montaruli, 5 - 7 Apr. 2005

[Greisen 66; MACRO,

Zatsepin & Kuzmin66]

2.73 K

Threshold for GZK cut-offThreshold pr p-gDpN

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

Teresa Montaruli, 5 - 7 Apr. 2005

Integrating over Planck spectrum Ep,th~ 5 ·1019 eV

[Greisen 66; MACRO,

Zatsepin & Kuzmin66]

GZK cut-off?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 MACRO,

Galaxy cannot contain EHECR: at 1019 eV Larmour radius of CR p comparable to Galaxy scale

AGASA: E>4 ·1019 eVno evidence of anisotropies due to galactic disc butlarge 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

Berezinsky et al, 1985 MACRO,

Gaisser, Stanev, 1985

Neutrino production: bottom upBeam-dump model:p0 g-astronomy p± n-astronomy

Neglecting g absorption (uncertain) n g

Targets: p or ambient g

Teresa Montaruli, 5 - 7 Apr. 2005

From photon fluxes to MACRO, n predictions:pp

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

Teresa Montaruli, 5 - 7 Apr. 2005

Exercises!

From photon fluxes to MACRO, 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 0.1·Ep /2

Teresa Montaruli, 5 - 7 Apr. 2005

Magnetic MACRO,

Cloud

Magnetic

inhomogenities

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

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

Teresa Montaruli, 5 - 7 Apr. 2005

CasA Supernova Remnant in X-rays MACRO,

Shock fronts

Magnetic

Inhomogenities

v2>v1

John Hughes, Rutgers, NASA

Interstellar medium

Blast shock

1st order Fermi accelerationThe 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: r1v1=r2v2

For ionized gas r1/r2 = 4 v2 =4 v1

upstream

v1=u/4<v2

v2=u

Shock front at rest: upstream gas flows into shock

at v2 =u

And leaves the shock with v1 = u/4

downstream

Teresa Montaruli, 5 - 7 Apr. 2005

Particles lost in each round trip on the shock MACRO,

Fermi mechanism and power lawsE = 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

Teresa Montaruli, 5 - 7 Apr. 2005

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