Progresses in Neutrinos Astronomy -...

Progresses in Neutrinos Astronomy

Francesco VissaniINFN, Gran Sasso Theory Group

Neutrino astronomy is a discipline aimed at investigating a class of high energy

phenomena, by exploiting certain specific nuclear reactions. It enjoys a steadily

improving instrumentation and observational successes, whereas theoretical

investigations often suffer of an artificial separations among astrophysics, nuclear and

particle physics. We review recent progresses focussing in particular

on the sub MeV–100 MeV energy region.

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.

(2 very important)NEUTRINO

INTERACTIONS

F. Vissani Rome, November 15, 2010

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Elastic scattering (ES) reaction

Figure 1: Due to the ES reaction, the (anti) neutrinos of several MeV hit the atomic

electrons and produce events in the forward direction. We can distinguish ES events if

we know the direction of the source.

A sample of electrons from ES events locates its source, even if not

identified by more conventional astronomical means.


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Figure 2: Image of the Sun obtained by exploing the electrons from the ES reaction

as obtained by the Super-Kamiokande neutrino telescope, to be discussed in a while.

Similar results first obtained by Kamiokande and the detection of SN1987A are the reasons of 2002 Nobel prize awarded to Koshiba.


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Inverse beta decay (IBD) reaction

Figure 3: IBD: an almost isotropic reaction occurring when Eν > 1.8MeV with a

much larger cross section: σIBD ∼ G2FmpEν whereas σES ∼ G2

FmeEν . This reaction

is important, for free protons abundant in H2O and CnHm targets.

Positron is the footprint of the reaction; occasionally, even the

neutron can be observed, providing a double tag.


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-2 -1 1 2

-2

-1

1

2

10 20 30 40-1

-0.5

0

0.5

1

Figure 4: Left panel: direction of the events from a future galactic supernova; ES

events are mostly in the center, IBD events are instead almost isotropical. Right panel:

expected energy-cosine distribution of the events; energy is in MeV.

...of course these are simulated events!!!


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(some modern & active)NEUTRINO

TELESCOPES


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Figure 5: The Super-Kamiokande detector, located in the Kamioka mine (Japan). It

reveals ultrarelativistic charged particles thanks to their Cerenkov-Vavilov radiation.

Observes the electrons/positrons produced by neutrinos/antineutrinos

above ∼ 5 MeV. Very high statistics, thanks to its 50 kton mass of H2O.


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Figure 6: The Borexino detector, located in the Gran Sasso laboratory (Italy). Its

ultrapure scintillator permits us to observe the electrons-or positrons-which are occa-

sionally hitted by neutrinos-or antineutrinos, even when their energy is below MeV.

Observes another signature of νe: the neutron that follows from IBD and

that reacts with free protons by n+ p→ D + γ(2.2 MeV).


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Figure 7: The IceCUBE detector (South Pole). It reveals the (anti) muons occasion-

ally produced by (anti) neutrinos. It can get some information on supernova (∼ 10

MeV) νe’s thanks to its huge mass but it is optimized for high energy (>TeV) neutrinos.


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SOLARNEUTRINOS


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.

2pd+e++νe 2p+ed+νe

d+p3He+γ

3He+3He4He+2p

7Be+e‐7Li+νe 7Be+p8B+γ

3He+4He7Be+γ 3He+p4He+e++νe

7Li+p24He 8B24He+e++νe

(pp) (pep)

(hep)

(B)

(Be)

LacatenaPP

PPI

PPII PPIII

Figure 8: The main chain of nuclear energy generation in the Sun.

Directly tested by the relatively rare but energetic νe from PPIII branch (Super-Kamiokande,

SNO, Borexino) and by the more important PPII (Borexino). The PPI neutrinos are determined

indirectly, i.e., from luminosity constraint, if we assume the Sun in equilibrium state.


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Recent measurements of the PPIII neutrinos

Figure 9: Recent measurements of Boron-8 neutrinos (PPIII branch); ‘salt’ and ‘prop.

counter’ measure the neutron from ν D → ν p n reaction [i.e., deuterium dissociation

by neutral currents]. SK=Super-Kamiokande; BX=Borexino. The last points are new.


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Energy dependence of solar neutrino oscillations

Figure 10: The solar neutrino conversion probability [Ianni 2010]

The downturn at high energies is neatly explained by matter (MSW)

effect, i.e., coherent scattering of electron neutrinos on solar electrons,

which is not relevant for atmospheric neutrino oscillations instead.


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Recent measurements of photospheric heavy element abundances

lead to solar models unable to reproduce the helioseismic results.

This “solar composition problem” demands a re-analysis of inputs and

assumptions; opacity, distribution of the metals, tests of secundary chain

(i.e., CNO neutrinos); role of helioseismological constraints.

..............

This can be explored efficiently approximating in linear response the

output of evolutionary solar codes: the linear solar model [Villante Ricci 2009].

E.g.; this new tool has been recently used to argue against speculations on

“dark baryons” in the core of the sun [Villante 2010].


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ANTINEUTRINOSfrom

THE EARTH


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νe from terrestrial radioactivity

232Th and 238U beta-decay chains give rise to observable νe [Eder 66; Marx 69].

The first type of nuclei are 4 times more abundant; positrons produced in

the IBD reactions extend to ∼ 1.5 MeV and ∼ 2.5 MeV, respectively.

**************

A hint of observation is due to KamLAND in 2005, 2008 in noisy

conditions and concurrent νe signal resulting from nuclear reactors.

Few months ago, the first high confidence detection

in a clean detector was obtained by Borexino, putting

geo-neutrino science on firm observational bases.


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Borexino data

best-fit

eνreactors

eνcontribution from geo-

background

Light yield of prompt positron event [p.e.]500 1000 1500 2000 2500 3000 3500

Eve

nts

/240

p.e

./252

.6to

n-y

ear

0

1

2

3

4

5

6

7

8

Figure 11: The νe signal is due in equal proportions to νe’s from far away reactors and

those from terrestrial radioactivity: The latter component dominates at low energies.

By future data, 232Th and 238U components should be distinguishable.


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SUPERNOVA(anti)NEUTRINOS


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Appointment with the next galactic supernova

0. 1. 2. 3. 4. 5.t@secD

0.6

2.

4.

6.

LΝ

e@1052ergsecD

0.00 0.05 0.10 0.15 0.20t@secD

10

20

30

40

50Evis@MeVD

Figure 12: Left: luminosity of νe as deduced from SN1987A observations. Right:

simulated events in Super-Kamiokande, assuming a supernova located at 10 kpc.

Using the times of the neutrino events (IBD+ES) and their directions

(ES), one can infer with a precision of at least ∼ 10 ms the time of the

burst, that is comparable with its expected duration [Pagliaroli et al., PRL 2009].

A precious information for Virgo and LIGO!


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One practical problem with supernovae

Supernovae are very important and interesting for astronomy, nuclear

physics and particle physics. SN1987A shows that we can learn a lot

studying their neutrinos.

But we have to face the following issue:

Galactic core collapse events are rare for human

standards and neutrino experimentalists are

definitively human.

How could we possibly overcome this limitation?


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The diffuse supernova neutrino background

From the cosmic density rate of core collapse supernovae, RSN, and the

average energy spectrum of νe, dn/dE, we calculate the diffuse flux:

dNνe

dt da dE=

c

H0

∫ zmax

0

dzRSN(z)√

ΩΛ + Ωm(1 + z)3× dn

dE(E(1 + z))

[Zel’dovich & Guseinov 65; Ruderman 65]

This is a steady signal: we do not need to wait our life to collect it!

Note the cosmological redshift of the energy spectrum.

Let us discuss the present chances to detect it.


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1) there is a window for the experimental search

20 30 40 50 60 70 80Visible Energy Ee [MeV]

0

0.5

1

1.5

2

2.5

3

3.5

dN / d

E e [ (22

.5 kto

n yr)-1 M

eV-1 ]

Figure 13: 2003 spectrum of Super-Kamiokande. The maximum relic supernova νe

signal is shown in light blue. Energy threshold set at 19.3 MeV [Malek et al., 2003]

Causes of background: atmospheric neutrinos at high energies; muons

below the thereshold for light at low energies. If the IBD neutron could

be seen, background could be decreased & energy threshold lowered.


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2) cosmic rate of supernovae is to some extent known

10 100

10-4

10-3

R CCSN

[yr-1 M

pc-3 ] Kistler et al. (2008)

Smartt et al. (2008)Dahlen et al. (2008)

Cappellaro et al. (1999)Dahlen et al. (2004)Cappellaro et al. (2005)Botticella et al. (2008)

10 100distance D [Mpc]

0.1

1

Ia/CC

ratio

0.5 1redshift z

99%90%

Figure 14: The cosmic density rate of core collapse supernovae RSN(z) obtained

rescaling the star formation rate (light brown points) does not disagree much with

direct supernova observations [according to Beacom 2010].


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3) the obstruction for a neat theoretical prediction

We ignore the high energy emission of the average supernova. Lacking a

definitive theoretical model, we explore the consequences of the assumption

that SN1987A νe emission is the typical outcome.

15 20 25 30 35 40EΝ in MeV

0.05

0.10

0.15

0.20

0.25

FHEΝL in eventsHMeV*yr*22.5 ktonL

0.2 0.4 0.6 0.8 1.0 1.2 1.4z

0.2

0.4

0.6

0.8

1.0

1.2

1.4

FHzΝL in eventsHyr*22.5 ktonL

Figure 15: Differential distributions of the event rate in 22.5 kton (Super-Kamiokande

fiducial volume). Left panel: Events per MeV per year. Right panel: distribution in

redshift, for the events above 19.3 MeV, used in Super-Kamiokande 2003.


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10 15 20 25 30 35 40

0.050.10

0.501.00

5.0010.00

Eth @MeVD

N ev@y

r-1 D

Cumulative Distribution in Super-Kamiokande

Figure 16: The 3 lines correspond to the uncertainty in the cosmic rate of SN explo-

sion. The error-bars are the 2 sigma ranges due to the uncertainties in the astrophysics

of the explosion, as evaluated using SN1987A observations [Pagliaroli & FV 2010].

If we can lower the threshold, the rate increases and the error decreases.


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HIGH ENERGYNEUTRINO

DETECTORS


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Theoretical expectations

Certain astronomical objects are thought to be ‘point sources’ of very high

energy ν, whose detection is the dream of modern neutrino astronomy.

A principal example is provided by supernova remnants:

Figure 17: Left: a supernova remnants that accelerates CR’s, subsequently

interacting with a molecular cloud [Drury, Volk, Aharonian 94]. Right: The principle of

detection of νµ: the induced, upgoing muon flux [Zheleitnik 57; Markov 60, Greisen 60].


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Mesons yield neutrinos, e.g., π± → µ±(−)νµ but we have also π0 → γγ!

Thus, when VHE gamma’s are measured, we get an upper bound on neutrinos.

1 5 10 50 100 500 1000EΓ in TeV10-20

10-18

10-16

10-14

10-12

IΓ in 1Hcm2 s TeVL

Γ-ray sources with 1 Ν

per km2´year above 1 TeV

γ-ray sources more intense than 2× 10−15/(cm2 s TeV)

at 20 TeV are particularly interesting for neutrino search[Aharonian, Sahakyan, FV, 2010]


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There are a few SNR’s that satisfy this conditions!!!

-1 -0.5 0 0.5 1 1.5 2 2.5Log@EΝ1 TeVD

-12

-11.8

-11.6

-11.4

-11.2

-11

Log@FΝ

EΝ

2.5

1cm-

2s-

1T

eV1.

5D

-1 -0.5 0 0.5 1 1.5 2 2.5Log@EΝ1 TeVD

-12

-11.8

-11.6

-11.4

-11.2

-11

Log@FΝ

EΝ

2.5

1cm-

2s-

1T

eV1.

5D

For a detector of 1 km2, with threshold set at 1 TeV,

we expect up to 2.5 signal events per year from RX

J1713-7-3946 and 1 background event [Villante, FV 08].


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An overview of the experimental progresses

[Compiled by IceCUBE collaboration]


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• Neutrino astronomy is a stimulating field, often considered a

branch of particle physics, but belonging with the same / or

more / rights to astronomy.

• Observations of low energy νs in very good shape, thanks to

solar νs and to a new branch: geo-neutrinos.

• Theory of solar and supernova νs also proceeds, though slowly.

Speculations and overoptimistic statements are common in / low

and high energy / neutrino astronomy.

• High energy ν telescopes are a reality. Till now, no source

seen yet, but observational perspects are becoming more

definite. There is space for progresses and perhaps surprises.


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