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ASTROPARTICLE PHYSICS
NEUTRINOS
COSMIC RAYS
DARK MATTER
CMB
BARYOGENESIS
COSMOLOGYASTROPHYSICS
PARTICLE PHYSICS
Esteban Roulet, Bariloche, ICFA2010
GAMMA RAYS
DARK ENERGY
Some basics of cosmology:
Cosmological Principle: isotropy + homogeneity => FRW metric
ds2=dt 2
−a2 dr 2
1−kr 2r 2d2 k=1 closedk=1 closedk=0 flatk=−1 open
Einstein's equation: aa 2
k
a2=8G
3
Scale factor
Expansion rate (Hubble 'constant'): H=aa
Critical density: c≡3 H 2
8Gi≡
i
c
Observations consistent with flat Universe (k=0, i.e. )
Ho= 71 +/- 4 km/s/Mpc age of the Universe 13.7 +/- 0.2 Gyr
=mr
total=1.02±0.02
Equation of state : P=w w=0matter , 1/ 3rad ,−1
a t
⇒∑i=1k
a2 H 2
aa=−
4G3
3P
w−13⇒( acceleration)
Evidences for dark matter
Flat rotation curves in galaxies
X-ray T profiles in clusters → gravitational potential Gravitational lensing in clusters
Redshift zM
ag
nit
ud
e m
(0,1)(.5,.5)(1,0)(1.5,-.5)
(ΩM, ΩΛ) = (0,0)(1,0)(2,0)
Stars
BaryonicDM
DarkEnergy
Non baryonic
DM
TYPE I SN
(m,
)
redshift
Lu
min
osi
ty
accelerated expansion
CMB T anisotropies
WMAP
PRECISION COSMOLOGY
The universe is expanding → it was smaller and hotter in the past
The early Universe was ionized, and hence opaque to photons until recombination: p+e → H that took place when T ~ 0.3 eV, i.e.for redshift z = 1100
those photons arrive to us as microwaves → CMB, and the density inhomogeneities present at decoupling lead to T anisotropiesat present that depend on the cosmological parameters:
n≃400 cm−3
Power spectrumof T fluctuationsvs. angular scale
o
e≡1z
SOME CMB EXPERIMENTS:
WMAP BOOMERANG DASI
PLANCK
Taking data, results by 2012
COBE
PLANCKsimulation
TYPE IA SN SEARCHES:
Measure acceleration and eq. of state of dark energy
cosmological constant: w = -1, quintessence: w > -1 and t dependent
Present::
SNLS (Supernova Legacy Survey at CFHT)+ ESSENCE at CTIO + SCP + ....Union 2008 results: w = -0.969 +- 0.069 +- 0.063new results expected soon
Future: DES (Dark Energy Survey at Cerro Tololo)
SNAP (Supernova Acceleration Probe by 2020?)
w=P
Going farther back in time, when T~ 0.1 – 1 MeV → PRIMORDIAL NUCLEOSYNTHESIS
GF2 T 5
~G~ g T4
M Pl2
n in equilibrium until they freeze out when = H ->
equilibrium : pe⇔n
# dof ( i )
T fo≃0.8MeV ⇒np≃exp −
mn−mp
T fo
≃16
Deuterium formed disintegrates until T < 0.1 MeV, when n/p~1/7 (due to some n decays)most neutrons end up in He nuclei →
He mass fraction:Y He=4n /2n p
≃0.24
Baryon density
Number of families < 4.3
More than luminous
but not enough to account for all DM
non-baryonic DM is also required for structure formation
At T~ 3 MeV, neutrinos decoupled → Cosmic background with
At even earlier times (T~ GeV), dark matter WIMPs (Weakly Interacting Massive Particles)may have decoupled:
h2≃
0.1 pb c
⟨ A v ⟩h≡
H 0
100km / s /Mpc≃0.7
G F2 m2
4≃0.1 pb m
100GeV 2
Note that:
Stable and neutral (weakly interacting) massive (~100 GeV) particles areperfect candidates to contribute sizeably to
most appealing WIMP candidate: the supersymmetric neutralino
Decouple when T~m/20
T =1.9oK ,ni
≃100
cm3
SUPERSYMMETRY
Particles superpartners
fermions sfermionsgauge bosons gauginos charginosHiggses (H
1 , H
2) higgsinos neutralinos
graviton gravitinos
=c1 c2 Zc3h10c4
h20
Neutralino annihilations:
Weak interactions → easy to get significant relic density
The neutralino is usually the lightest susy particle (LSP), stable due to R parity
DIRECT WIMP DETECTION: measure nuclear recoil of halo WIMP interacting in detector
E rec12m v
2~50keV
m
100GeV ⟨ v300 km / s ⟩
2
Rates depend on scattering cross section:
spin-independent:: ~ A^2 enhancement for heavy targets
spin-dependent:: ~ J(J+1) not enhanced
Need low thresholds and tiny background(underground labs, ultrapure materials, ...)
mqq qQuark scalar coupling ~ contribution from strange coupling to gluons relevant
CDMS (CRYOGENIC DM SEARCH at SOUDAN)19 Ge (230 g) and 11 Si (100 g) detectors at 50 mK
measure phonons and ionization to distinguish nuclear recoils from gamma backg. (e recoils)dec 2009: 2 events (p=23% of being backg)
DAMA (250 kg NaI scintillators at Gran Sasso)
Observe annual modulation due to Earthmotion relative to DM distribution
Spin indep
SUSY
Threshold statistics
(Gelmini et al)
SUPERKAMIOKANDE
10−10 pb
SupeKamiokande (50 kton water Cherenkov)looked for upgoing muons from Sun's direction:
Bound relevant for low masses and for spin-dependent interactions (Sun made mostly of H)
Spin independent Spin dependent
INDIRECT DETECTION:look for GeV neutrinos from WIMP annihilations in the sun or earth
( annihilation rate = Capture rate / 2 ) A
R
A T=N A∫E th
m
dE
d
dE
∫E th
E
dE
d
dE
Range
INDIRECT WIMP DETECTION II:
look for GeV gammas, positrons or antiprotons from WIMP annihilations (or decays) in the halo (hot topic in recent past)
example:Gamma rays positrons antiprotons
background 200 GeV Wimp signal ?
Experiments:baloons with spectrometers, calorimeters,.. (PAMELA, ATIC,...)GeV gamma ray/electron detectors in satellites (EGRET, GLAST,...)TeV gamma Imaging Air Cherenkov Telescopes (HESS, MAGIC,...)
recent results:
GeV gammas
EGRET GeV gamma excess not confirmed by Fermi GLAST, results agree with background
GLASTSi tracker
electrons:
ATIC excess not confirmed
Fermi and HESS harder than conventional diffusion model
Need extra component:WIMP annihilations? or pair production in SN remnantsor in nearby pulsars?
positrons:
PAMELA 'excess'
backg ?
Antiprotons:
pre-Pamela
Pamela
No indications requiring components beyond standard CR spallation
Additional WIMP searches:gammas from annihilations in the galactic center,annihilations in dwarf satellite gallaxies, monochromatic gamma lines from one-loop annihilations,
-> no clear signals yet
THE NEUTRALINO DETECTION
q q
An
nih
ilation
(ind
irec
t dete
ctio
n)
Scattering(direct detection
Pro
du
ctio
n (
LH
C)
SUSY discovery at LHC may be the first strong indicatication in favor of WIMPs
e.g. trilepton signal with missing ET :
(but this will not prove that they are the DM)
p p20 l1
0l l−10
A BIT OF NEUTRINO ASTROPHYSICS
Neutrino underground detectors first built to study proton decay,the background was from atmospheric neutrinos, and the threshold was moved down below 10 MeV to detect solar neutrinos, justin time to detect SN1987A
Massive star (~20 Msun)burned H -> He-> C, O -> Si -> Feand exploded when Fe core reached Chandrasekar mass
hot colapsed core cooled byneutrino emission
10-40 MeV neutrinosobserved during ~ 10 sec
with SuperKamioka, a galacticSN will give thousands of events(but 2/century)
~3% of solar energy emitted in νe with E
ν < 14 MeV
4p 4 He2e2e27 MeV
SOLAR NEUTRINOSSolar Energy produced by fusion:
EC
16 d71Gaee−
30 ton
Gallex / GNO Super Kamiokande
Cherenkov in H2O
Exp = 0.3÷ 0.6 Theory
ee−ee−
71Gae e71Ge
Atmospheric Neutrinos
oscillations νµ ⇔ ν
τ
R e≡ee
≃2⇒
But observations: Rµ e
≈ 1 and angular dependence! (SuperK '96)
m2≃2.5×10−3eV 2 ; sin 2≃1
Losc
≈ 103 km E[GeV]
Cosmic ray (p, He,...,Fe)+air → π , K + ...
ee
Masses and mixings further constrained by reactor and accelerator long baseline expts.
Future experiments :
measure θ13
102÷ 103 (reactor, LBL)sign ∆ m2
atm (LBL, SNe)
masses (Katrin,Gerda,Cuore: 1→0.1eV cosmology: Σm
ν <0.7 → 0.1 eV)
CP violation (µ Factory) + β beams, off axis LBL, ...
Lot of ideas and hard work devoted
∣matm2 ∣=2.40±0.1210−3eV 2
Oscillations among 3 neutrinos
sin 2atm=0.50±0.07
m sol2=7.67±0.2310−5 eV 2
sin 2sol=0.304±0.022
sin 213≤0.056
Yukawa coupling: λ νL ν
R H0 ?
How to give mass to the neutrinos ?
⇒ mD≃1 0 −1 eV
1 0 −1 2 ≤1 0 −1 2
But why ? , why ? m≪me
SeeSaw mechanism (Yanagida; GellMann, Ramond, Slansky '79)
ν is light because νR is heavy:
Lm∝L ,R 0 mD
mD M L
R
⇒ m≃mD2
M≃10−1 eV mD
GeV 2
1012GeVM
νR
νL
⇑
m
_
⇑⟨H 0
⟩≃245GeV
Leptogenesis (Sakharov, FukugitaYanagida)
Why there is more matter than antimatter, with ?
After the Big Bang, when T > M, νR was in equilibrium
⇒ n(νR) ≈ n(γ ) ∝ T3
for T < M, n(νR) ∝ exp(T/M)
⇒ νR tries to decay, but does it at a slower rate than the
expansion of the Universe (⇒ out of equilibrium)
If CP is violated:
⇒ a matter – antimatter asymmetry is generated
note that:
R LH ≠R L H ∗
m≃10−10mp , n≃109nB⇒≃0.1B
nB−n B
n≃10−9
THE COSMIC PIE AGAIN:
LECTURE 2
ASTROPARTICLE PHYSICSEsteban Roulet, Bariloche, ICFA2010
THE ENERGETIC UNIVERSE
rays (Fermi) (Amanda)
UHE Cosmic rays (Auger)Multimessenger astronomy
p
Auger
HiRes
Pulsar
GRB
AGN
SNR
Radio Galaxy
Examples of Astrophysical Objects
Colliding galaxies Diffuse
emission
max≃3mm
T / oK
Production mechanisms for photons
Thermal:
e+e- annihilations:
Bremsstrahlung:
Synchrotron:
Inverse Compton:
Hadronic production:
e+
e-
Interstellar matter
e-
hot surface
magnetic field
e-
B
e-
pInterstellar matter
e-
e-
e-
Diffuse gammas from Galaxy
Multiwavelength astronomy (from radio to 100 GeV)
atmosphericattenuation
Radio telescopes
COBE/WMAP/Planck
optical telescopes
INTEGRAL/SWIFT/...
CGRO/Fermi
TeV photon detection:Imaging Atmospheric Cherenkov Telescopes (Whipple, Veritas, Magic, Hess)
EM cascade: (e+ e- gamma)
e+- produce Cherenkov light if
≡n−1≃3 10−4exp −z /7.3km E th=thme≃40 MeV at z=10 km
Cherenkov angle
Cherenkov light spread over ~ 200 m at ground
C=cos−1 1n ≃0.7oat 10 km
≡vc
1n
Note that in water, n=1.3
E the ≃0.26 MeVc≃41o
TeV Astrophysical Sources
HESS : Dark MatterGamma Ray spectrum from Galactic CentreGamma Ray spectrum from Galactic Centre
Data follows power law typical of astrophysical sources
z=0.165 BLLac (H2356-309 )
distant sources strongly attenuated by background photons (starlight, CMB, radio,,..):
Can measure IR background from observed attenuation
Sync
IC
TeV
e e−
IC
brems
0
Discriminating leptonic vs. hadronic scenarios in CasA
(a way to know if protons areindeed accelerated in SNR)
High energy emission (γ -ray):- self-compton (electro-magnetic) ?- π 0 decay (hadronic) ?
π + −
ν µ µ + −
ν µ ν e e+ −
Low energy emission (X-ray) :Synchrotron emission of e- in jet
e-
γ
γ p
π ογγ
e-
γ
νν can help to distinguishcan help to distinguishhadronic and leptonichadronic and leptonic
accelerationacceleration
Neutrinos produced in beam dumps fromaccelerated protons
ANTARES NEMO NESTOR
NEUTRINO TELESCOPES (10 GeV to PeV)
Amanda
km3 detector at South Pole, looking at northern sky (complete by 2011)
km3 detector at Mediterraneanlooking at southern neutrino sky (proposed km3NET)
One may even distinguish neutrino flavors
muon neutrino
electron neutrino
tau neutrino (double bang)
Atmospheric neutrinos measured up to 100 TeV, no point sources yet
High energy atmospheric neutrinos
Atmospheric ν s mainly from pion decays at low energies,
but above 100 GeV pions are stopped before decay =>kaons become the main source,
but above ~1014 eV prompt charm decays dominate
L= cdecay length
L≃6 kmE
/100 GeV
LK≃7.5 kmE K /TeV
LD≃2 kmE D/10 PeV π
K
c
Prompt charm productionsample gluon density distribution at=> x
2 < 10 5 for E>1015 eV
need to extrapolate from measured valuesalso requires to include NLO processes
x 2≃M cc
2
2xF s
> one may learn about particle physics!
Co
smic
ra
y f
lux
Energy
CHARGED COSMIC RAYS
E3 x FLUX
knee 2nd knee ankle GZK ?
Kascade
SCIENTIFIC OBJETIVES:
Spectrum: understand the origin of the different spectral features
Arrival directions: search for anisotropies (identify the sources)
Composition: light or heavy nuclei, photons, neutrinos, others?
Study of interactions at energies unreachable at accelerators
PeV energies
MILAGRO water Cherenkov
ARGO in Tibet 6500 m2 of RPC
1 km2
present and future sensitivities
1600 detectors instrumenting 3000 km2 and 24 telescopes
THE PIERRE AUGER OBSERVATORY
the Auger Collaboration: 18 countries, ~400 scientists
at the highest energies, only few cosmic rays (CR) arrive per km2 per century !to see some, a huge detector is required, e.g.
surface detector
fluorescence detector
Previous experiments:
AGASA: (Akeno, Japan 1990-2004) Area: 100 km2
111 Scintillators ( e+e- ) and 27 shielded proportional counters (muons)
Fly's Eye (1981-1993) Utah, USAHiRes (1997-2006)
Fluorescence telescopes
Also Volcano Ranch, Haverah Park, Yakutsk, Sugar, ...
the GreisenZatsepinKuzmin effect (1966)
protons cannot arrive with E > 6x1019 eV from D > 200 Mpc
AT THE HIGHEST ENERGIES, PROTONS LOOSE ENERGY BY INTERACTIONS WITH THE CMB BACKGROUND
Feγ = For Fe nuclei:
after ~ 200 Mpc the leadingfragment has E < 6x1019 eV
ligther nuclei get disintegrated on shorter distances
1 Mpc 100 Mpc
pγ π o ppγ πn
⁰( produce GZK photons)
±( produce GZK neutrinos)
A A 'nucleons
THE END OF THE SPECTRUM: GZK?previous results
AGASA: NO HiRes: YES
(ICRC09)
AUGER and HiRes SPECTRA
Some basics on air showers:
ELECTROMAGNETIC SHOWERS ( e+ , e- , )
Energy loss of electrons:
Critical E: ionization loss = loss by particle production (brems)
Heitler model: after , particles split in twoafter n generations, number of particles
Shower growth stops when
and depth of maximum is
dEdX
=−E −EX R
Ec=X R ⟨ E⟩≃86 MeV
em=ln2 X RN=2n , with n=X /em
E0/N=E c ,i.e. N max=E0/E c ≃1011 E0
1019 eV X max=n em=X R ln E0/Ec
for air X R=37g
cm2
X
HADRONIC SHOWERS
each interaction produces pions (multiplicity)
Number of generations from:
Energy of em component:
Estimating as the maximum of the first generation s:
For nuclei: nucleons with
n tot
nch=2 n tot /3 ± E Edec
E0/n totn=Edec typically n~5−6
nneut=n tot /3 0 2
Eem=E0−2/ 3n E0 ~0.9 E0 for 1019 eV
X max
En=E0/ A
0
X max=IX R ln E 0/ntotEc
depends on and on multiplicity
I~ p−air−1
em component
reinteract until
A
CROSS SECTIONS MULTIPLICITIES
Energy flow at LHC
=−ln tan/2
COMPOSITION FROM Xmax
CONFLICTING RESULTS ?
XmaxRMS Xmax
AUGER SD photon bound
photon showers mostly electromagnetic: long rise times
and develop later → small curvature radius(and large Xmax in FD)
photon fraction:
< 2% at E > 10 EeV
< 31% at E > 40 EeV excludes most top-down models
Astroparticle Physics 29 (2008) 243–256
GZK
AUGER BOUNDS ON DIFFUSE NEUTRINO FLUX
unlike hadronic CRs, neutrinos can produce young horizontal showers above the detector, and upcoming near horizontal tau lepton induced showers
young (em) shower
old (muonic) shower
Horizontal young showers?60% of tank signals with large Area / peakElongated tracks: L/W > 5 Propagation with v ~ c
ZERO CANDIDATES
( E-2 )
p
He
O
Fe
Hooper, Sarkar, Taylor astro/0407618
If GZK neutrinos were observed,it would be a strong hint favoring a light composition
only for E/Z ≫ 1019 eV deflections in galactic magnetic fields become less than a few degrees and CR astronomy could become feasible
COSMIC RAY ASTRONOMY ?
≃3o ZB
3GLkpc
6×1019 eVE
galaxies up to 100 Mpc
the nearby Universe is quite inhomogeneous
the radio sky
CenA
M87
Crab
Geminga
VelaCygnus Sagitarius
CasA
supernovae: preferred candidatesources for E < 1018 eV
active galaxies: plausible candidates for E > 1018 eV
Auger 2007: from 27 events with E > 57 EeV, 20 are at less than ~3 degrees from an active galaxy at less than ~ 75 Mpc , while only 6 were expected
* nearby active galaxies
CR
CenA
Excess around Centaurus A
Closest Active galaxy: Centaurus A
2 events at less than 3 deg from it
within 18 deg:12 observed/2.7 expected
central black hole withmore than 100 million solar masses !
collision of 2 galaxies
relativistic jet
(~ 3 Mpc)
HESS observation of Centaurus A (0.1 – 10 TeV gammas)
COSMOLOGY MARCHES ON ...COSMOLOGY MARCHES ON ...
the questions are similarthe tools are sharper
the answers are closer
