progresses in neutrinos astronomy -...
TRANSCRIPT
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
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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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