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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