the missing 95%: theory and phenomenology of …home.fnal.gov/~jkopp/pdf/dmde.pdf(inferred from...

The missing 95%: Theory and Phenomenology ofDark Matter and Dark Energy

Joachim Kopp

DPG Spring Meeting Göttingen, March 2012

Fermilab

Joachim Kopp Dark Matter and Dark Energy 1

Outline

1 Evidence for dark matter

2 Finding dark matterDirect detectionIndirect detectionProduction at colliders

3 Modelling dark matter

4 Dark energy


Outline




4 Dark energy


Celestial mechanics

Stellar/galactic dynamics relates:

The mass distribution(inferred from brightness)

Kinetic energy(inferred from Doppler shifts) Fritz Zwicky

1898–1974Vera Rubin

1928–

Observations of rotational velocities in galaxies show: Rubin 1975

The gravitational pull onperipheral stars is stronger thanpredicted from the mass of theluminous matter M

mv2

r= GN

mMr2 at r →∞


Collisions of galaxy clusters

Artist’s rendering (Image: NASA)

red = gas (from x-ray observations)blue = (dark) matter distribution (from gravitational lensing)


Collisions of galaxy clusters

Image: NASA (Chandra [x-ray], ESO WFI [lensing], HST [optical])

red = gas (from x-ray observations)blue = (dark) matter distribution (from gravitational lensing)


The Cosmic Microwave Background (CMB)

WMAP’s observation of the CMB: A fingerprint of the universeat t ' 300 000 yrs(when electrons and protons first combined to form atoms).

red = overdense, hot regions (0 . . . + 200 µK)blue = underdense, cold regions (−200 . . . 0 µK)


Image credit: NASA



More useful: The CMB fluctuation power spectrum


Image credit: NASA



red curve = theory predictionblack points = WMAP data


Image credit: NASA



too little DM (0.04ρc) right amount of DM (0.22ρc) too much DM (0.74ρc)


Image credit: NASA

What is this stuff?Modified laws of gravity?

I Hard to explain all observationsMACHOs (Massive Compact Halo Objects)?

I Planets, Brown dwarfs, neutron stars, . . .I Ruled out as dark matter in the mass range 0.6× 10−7M < M < 15M by

searches for gravitational microlensingI Searches for candidate objects yield too few of them

Hot (relativistic) Dark Matter (neutrinos or other relativistic particles)?I Cannot explain large scale structure of the universe

(hot dark matter would smoothen the galaxy distribution)

Cold or Warm Dark MatterI Axions

F Ultra-light, but non-relativistic due to non-thermal productionI Gravitinos

F Only gravitational couplings→ bad for direct/indirect/collider detectionI WIMPs (Weakly Interacting Massive Particle)

F New, heavy, stable particlesF Should have some non-gravitational interaction with SM particles for production

in the early universe


Outline




4 Dark energy


Direct Dark Matter detection

Idea: A WIMP (Weakly Interacting Massive Particle) can scatteron an atomic nucleus.

t

f

χ

f

χ

Strategy: Look for feeble nuclear recoil

Problem: Many background processes (radioactive decays, cosmic rays, . . . )can mimic the signal


Direct DM detection — The experimental challenge

background

suppression


Direct detection results

10 100 1000WIMP mass [GeV]

10-9

10-8

10-7

10-6

10-5

10-4

10-3

WIM

P-n

ucle

on c

ross s

ection [pb]

CRESST 1σ

CRESST 2σ

CRESST 2009EDELWEISS-IICDMS-IIXENON100DAMA chan.DAMACoGeNT

M2

M1

CRESST-II, arXiv:1109.0702

Assumptions here: Elastic DM scattering ∝ target mass (often realized in SUSY)Joachim Kopp Dark Matter and Dark Energy 12

Direct detection phenomenology of alternative models

Previous slide: Elastic dark matter (χ) scattering through scalar current[(qq)(χχ)] or vector current [(qγµq)(χγµχ)] assumed⇒ Cross section ∝ target massIn models with different coupling structure, the relative detectionefficiencies of different experimental technologies may be different


Direct detection phenomenology of alternative modelsSpin-dependent couplings

I E.g. coupling through axial vector current [(qγµγ5q)(χγµγ5χ)]I Cross section ∝ target spinI Cannot explain DAMA, CoGeNT, CRESST results

Inelastic dark matter Tucker-Smith Weiner hep-ph/0101138

I There may be two DM states χ and χ′ with m′χ = mχ + δ (δ ∼ 100 keV)

I Scattering χN → χ′N ⇒ heavy target nuclei kinematically preferredI Could explain CRESST, but not DAMA JK Schwetz Zupan 1110.2721

Leptophilic dark matter Bernabei et al. 0712.0562; Fox Poppitz arXiv:0811.0399; JK Niro Schwetz Zupan arXiv:0907.3159

Isospin-violating dark matter Feng Kumar Marfatia Sanford 1102.4331

. . .









. . .


30 40 5010-41

10-40

10-39

10-38

10-37

mΧ @GeVD

WIM

P-nu

cleo

ncr

oss

sect

ion

ΣSI

@cm2 D

iDM, ∆=90keV

Limits: 90%

Contours: 90%, 3Σ

v0 = 220 kms, vesc = 550 kms

Xenon

-100

KIM

S

CRESST-II

DAMA

CR

ESST

comm

. run








. . .









. . .

Conclusion: Hard to explain all data simultaneously


Direct detection uncertainties

Large uncertainty in local DM density

Large uncertainties in DM velocity distributionI Scattering rate depends strongly on DM velocityI DM streams?I Debris flow?

Predicting WIMP–nucleus cross sections is difficultI Models predict WIMP–quark cross sectionI Need to know quark content of the nucleonI Especially problematic for Higgs-mediated scattering:

coupling ∝ quark mass⇒ sea quarks dominateI Need to know nuclear form factor

especially difficult for spin-dependent scattering



Large uncertainty in local DM densityLarge uncertainties in DM velocity distribution

I Scattering rate depends strongly on DM velocityI DM streams?I Debris flow?

Kuhlen Lisanti Spergel arXiv:1202.0007, graphics courtesy of Mariangela Lisanti






Large uncertainty in local DM densityLarge uncertainties in DM velocity distribution

I Scattering rate depends strongly on DM velocityI DM streams?I Debris flow?





Outline




4 Dark energy


Indirect Dark Matter detection

Idea: WIMPs (Weakly Interacting Massive Particles) χ can annihilate(or decay) into Standard Model particles (f ) in an astrophysical environment.

t

f

χ

f

χ

Strategy: Look for annihilation products in cosmic rays

Problems:Many other sources of cosmic raysPropagation of charged particles in the galaxy poorly understood

Advantage:Many sources to look at


Indirect DM detection — The experimental challenge

look atmany sources


Indirect DM detection — Examples

γ-rays from dwarf galaxiesIdea:Look for anomalous γ-ray flux

Pro:Few stars ⇒ few backgrounds

Con:

Relatively low DM densityResults model-dependentLarge astrophysicaluncertainties

Other indirect DM searches:Cosmic anti-matter (e+, p, . . . ) PAMELA, Fermil-LAT, . . .

γ-rays from the galactic center Hooper et al.

High-energy neutrinos from the Sun IceCube, SuperKamiokande

. . .



γ-rays from dwarf galaxies

101 102 103

WIMP mass [GeV]

10-26

10-25

10-24

10-23

10-22

10-21

WIM

Pcr

oss

sect

ion

[cm3

/s]

Upper limits, Joint Likelihood of 10 dSphs

3*10−26

bb Channel

τ+ τ− Channel

µ + µ− Channel

W+W− Channel

Thermal relic cross section=> Correct relic abundance

Fermi-LAT 1108.3546see also Geringer-Sameth Koushiappas 1108.2914

Idea:Look for anomalous γ-ray flux


Con:





. . .



γ-rays from dwarf galaxies

101 102 103

WIMP mass [GeV]

10-26

10-25

10-24

10-23

10-22

10-21

WIM

Pcr

oss

sect

ion

[cm3

/s]


3*10−26

bb Channel

τ+ τ− Channel

µ + µ− Channel

W+W− Channel



Idea:Look for anomalous γ-ray flux


Con:





. . .Joachim Kopp Dark Matter and Dark Energy 19

Outline




4 Dark energy


Dark matter at colliders

make your

own needles!


Generic collider searches for dark matter

Idea:Produce WIMPs in collisions of Standard Model particles

t

f

χ

f

χ

WIMPs can recoil against a jet or a photon from initial state radiation

t

q

χ

q

g

χ

t

e−

χ

e+

γ

χ

Experimental signatures: Mono-jets + /ET and mono-photons + /E


LHC limits on DM–quark couplings Plots from Fox Harnik JK Tsai 1109.4398see also work by Rajaraman et al. 1108.1196

90% C.L.

10-1 100 101 102 10310-46

10-45

10-44

10-43

10-42

10-41

10-40

10-39

10-38

10-37

WIMP mass mΧ @GeVD

WIM

P-nu

cleo

ncr

oss

sect

ion

ΣN

@cm

2 D

ATLAS 7TeV, 1fb-1 VeryHighPt

Spin-independent

Solid : ObservedDashed : Expected

ΧΓΜ Χ qΓΜq

Αs Χ Χ GΜΝGΜΝ

CDMSXENON-10

XENON-100

DAMA Hq ± 33%L

CoGeNT CRESST

90% C.L.

10-1 100 101 102 103

10-41

10-40

10-39

10-38

10-37

10-36

10-35

10-34

WIMP mass mΧ @GeVD

WIM

P-nu

cleo

ncr

oss

sect

ion

ΣN

@cm

2 D

ATLAS 7TeV, 1fb-1 VeryHighPt

Spin-dependent

Solid : ObservedDashed : Expected

ΧΓΜΓ5 Χ qΓΜΓ5q

XENON-10PICASSO

COUPP

SIMPLE

DAMAHq ± 33%L

Assumptions here:I Effective field theory approach valid (limits may be better or worse if EFT not valid)I Equal coupling to all quark flavors

Extremely competitive limits forI Light dark matter (below direct detection threshold)I DM coupled to gluons (high gluon luminosity at the LHC)I Spin-dependent DM interactions (DD suffers from loss of coherence)


Model-dependent collider searches: SUSY-DM

Idea:In many models, DM is produced in the decay of heavy, stronglyinteracting particles (for instance squarks and gluinos in SUSY)

Experimental signature: something + missing energyExample: pp → (g → jZχ0)(q → jjWχ0)

Advantage: Very sensitiveProblem:Minor modifications to the modelmay drastically change the phenomenologyProblem (all collider searches):Collider can only find DM candidate(s)


Outline




4 Dark energy


Electroweak-scale DM? — The “WIMP Miracle”

In the early universe, DM is in chemical equilibrium with other particles.

As the temperature drops, DM begins to annihilate away: χχ→ f fWhen the annihilation rate Γ(χχ→ f f ) drops below the Hubble expansionrate H, annihilations cease⇒ DM abundance remains constant (“thermal freeze-out”)From this requirement, and from the observed DM abundance today,cosmology predicts the DM annihilation cross section

〈σv〉 ' 3× 10−26 cm3/s

Consider generic DM coupling:

L ⊃ g2

M2 (χχ)(f f )

For typical coupling g ∼ 0.1, suppression scale M ∼ 100 GeV, DM massmχ ∼ 100 GeV, this yields the right value for 〈σv〉Conclusion: If dark matter originates from electroweak-scalenew physics, it automatically has the right abundance

The Wimp Miracle



In the early universe, DM is in chemical equilibrium with other particles.As the temperature drops, DM begins to annihilate away: χχ→ f f

When the annihilation rate Γ(χχ→ f f ) drops below the Hubble expansionrate H, annihilations cease⇒ DM abundance remains constant (“thermal freeze-out”)From this requirement, and from the observed DM abundance today,cosmology predicts the DM annihilation cross section

〈σv〉 ' 3× 10−26 cm3/s


L ⊃ g2

M2 (χχ)(f f )


The Wimp Miracle



In the early universe, DM is in chemical equilibrium with other particles.As the temperature drops, DM begins to annihilate away: χχ→ f fWhen the annihilation rate Γ(χχ→ f f ) drops below the Hubble expansionrate H, annihilations cease⇒ DM abundance remains constant (“thermal freeze-out”)

From this requirement, and from the observed DM abundance today,cosmology predicts the DM annihilation cross section

〈σv〉 ' 3× 10−26 cm3/s


L ⊃ g2

M2 (χχ)(f f )


The Wimp Miracle



In the early universe, DM is in chemical equilibrium with other particles.As the temperature drops, DM begins to annihilate away: χχ→ f fWhen the annihilation rate Γ(χχ→ f f ) drops below the Hubble expansionrate H, annihilations cease⇒ DM abundance remains constant (“thermal freeze-out”)From this requirement, and from the observed DM abundance today,cosmology predicts the DM annihilation cross section

〈σv〉 ' 3× 10−26 cm3/s


L ⊃ g2

M2 (χχ)(f f )


The Wimp Miracle




〈σv〉 ' 3× 10−26 cm3/s


L ⊃ g2

M2 (χχ)(f f )


The Wimp Miracle


χ

χ

f

f

←→

χ

χ

f

f



〈σv〉 ' 3× 10−26 cm3/s


L ⊃ g2

M2 (χχ)(f f )

For typical coupling g ∼ 0.1, suppression scale M ∼ 100 GeV, DM massmχ ∼ 100 GeV, this yields the right value for 〈σv〉

Conclusion: If dark matter originates from electroweak-scalenew physics, it automatically has the right abundance

The Wimp Miracle


χ

χ

f

f

←→

χ

χ

f

f



〈σv〉 ' 3× 10−26 cm3/s


L ⊃ g2

M2 (χχ)(f f )


The Wimp MiracleJoachim Kopp Dark Matter and Dark Energy 26

χ

χ

f

f

←→

χ

χ

f

f

Relating the DM and baryon abundances

Motivation: The DM and baryon energy densities in the universeare similar

ΩDM ' 5 Ωb

(Ω = energy density as fraction of “critical density” for flat universe)

If the DM and baryon number densities are similar and

mDM ∼ 5mp–10mp ∼ 5–10 GeV ,

this is quite natural.This is precisely the mass range where the direct detection hints(DAMA, CoGeNT, CRESST) have been observed!Baryon density Ωb generated by yet unknown dynamics behind theparticle–antiparticle asymmetry of the universe(not by thermal freeze-out)Assume dark matter (χ) density is also determined by χ–χ asymmetry⇒ Asymmetric dark matter




ΩDM ' 5 Ωb



mDM ∼ 5mp–10mp ∼ 5–10 GeV ,

this is quite natural.

This is precisely the mass range where the direct detection hints(DAMA, CoGeNT, CRESST) have been observed!Baryon density Ωb generated by yet unknown dynamics behind theparticle–antiparticle asymmetry of the universe(not by thermal freeze-out)Assume dark matter (χ) density is also determined by χ–χ asymmetry⇒ Asymmetric dark matter




ΩDM ' 5 Ωb



mDM ∼ 5mp–10mp ∼ 5–10 GeV ,

this is quite natural.This is precisely the mass range where the direct detection hints(DAMA, CoGeNT, CRESST) have been observed!

Baryon density Ωb generated by yet unknown dynamics behind theparticle–antiparticle asymmetry of the universe(not by thermal freeze-out)Assume dark matter (χ) density is also determined by χ–χ asymmetry⇒ Asymmetric dark matter




ΩDM ' 5 Ωb



mDM ∼ 5mp–10mp ∼ 5–10 GeV ,

this is quite natural.This is precisely the mass range where the direct detection hints(DAMA, CoGeNT, CRESST) have been observed!Baryon density Ωb generated by yet unknown dynamics behind theparticle–antiparticle asymmetry of the universe(not by thermal freeze-out)

Assume dark matter (χ) density is also determined by χ–χ asymmetry⇒ Asymmetric dark matter




ΩDM ' 5 Ωb



mDM ∼ 5mp–10mp ∼ 5–10 GeV ,

this is quite natural.This is precisely the mass range where the direct detection hints(DAMA, CoGeNT, CRESST) have been observed!Baryon density Ωb generated by yet unknown dynamics behind theparticle–antiparticle asymmetry of the universe(not by thermal freeze-out)Assume dark matter (χ) density is also determined by χ–χ asymmetry⇒ Asymmetric dark matter


Models of asymmetric dark matter

Example 1 Kaplan Luty Zurek, arXiv:0901.4117

B − L asymmetry generated athigh T (e.g. via Leptogenesis)

Effective superfield operator

L ⊃ 1M

X 2LHu (*)

transfers B − L ↔ 2X , e.g. via

χ0

ν ν

ν

X

X

ν

X

X

Final X (DM number) asymmetrydepends on # of SM speciescontributing to (*) at freeze-out

Example 2 Buckley Randall 1009.0270Blennow et al. 1009.3159

Generate X asymmetry inhidden sectorTransfer to B − L asymmetry inthe SM sector

I via B − L violatinginteractions (e.g. (*))

I via sphaleron processes

Example 3 Davoudiasl et al. 1008.2399Gu Lindner Sarkar Zhang 1009.2690

New heavy particles decaypartly into DM, partly into SMparticlesB − L− X is conservedDM (X ) does not participate inSM sphaleron processes⇒ Asymmetry frozen in


Outline




4 Dark energy


Evidence for dark energy: Type Ia Supernovae

When a white dwarf accretes matter from acompanion star, it becomes unstableonce it reaches ∼ 1.4M

I Re-ignition of nuclear fusionI Thermonuclear explosion

Since the progenitor mass is always ∼ 1.4M,all Type Ia Supernovae are very similar

I Energy release precisely knownI SN Ia are standard candles

Measurement:I Apparent brightness→ distanceI Redshift→ velocity

Result:I Long ago (very distant SN Ia,

low brightness), the universe wasexpanding more slowlythan we thought!

I It must be acceleratingCMB and Large Scale Structureobservations confirm this


http://hyperphysics.phy-astr.gsu.edu/hbase/astro/univacc.html

What is accelerating the Universe?

A cosmological constant?I An ad-hoc addition to the Einstein equations

Rµν −12

gµνR αα = 8πG Tµν + gµνΛ

I Observations require Λ ∼ (10−12 GeV)4

I Extra source of energy with negative pressure

QFT vacuum energy?I A vacuum expectation value (vev) or condensate of a quantum field

behaves like a cosmological constantI Problem: All known condensates/vevs are way too large!

(We expect Λ ∼ M4Pl ∼ (1019 GeV)4)


What is accelerating the Universe? (cont’d)

Quintessence: A new, slowly rolling scalar fieldI Introduce new scalar field φ slowly rolling down its potential V (φ)I Lagrangian:

Lφ =12

∂µφ∂µφ− V (φ)

I Energy and pressure:

ρ =12

φ2 + V (φ) , p =12

φ2 − V (φ)

I A cosmological constant corresponds to ρ = −p ⇒ require φ2 V (φ)

Extensions of general relativityI Scalar-tensor gravity: Modified Einstein-Hilbert action

S =1

16πG

Z p−g d4x R → S =

116πG

Z p−g d4x f (φ)× R

I A special case: f (R) gravity:

S =1

16πG

Z p−g d4x f (R)


for a review see Caldwell Kamionkowski 0903.0866

Summary

Overwhelming evidence for dark matterA lot of data available

I Direct detectionF Difficult to reconcile possible evidence with null results

I Indirect searchesF Strong exclusion limitsF Suffers from poorly understood astrophysical backgrounds

I Collider searchesF Generic searches (monojets + /ET , mono-γ + /E) and model-specific searches

(cascade decays) are underway full-steam

Dark matter modelsI Dark matter from electroweak scale new physics:

Correct cosmic abundance due to WIMP MiracleI Light (10 GeV) dark matter:

Correct cosmic abundance if related to baryon–antibaryon asymmetry

Dark energyI Accelerated expansion of the Universe well-establishedI So far, a cosmological constant is the leading explanation


Thank you!

Bonus material

Spin-dependent DM couplings?

Previous slide: Dark matter (χ) couplings through scalar current[(qq)(χχ)] or vector current [(qγµq)(χγµχ)] assumed⇒ Cross section ∝ target massAlternative: Axial vector [(qγµγ5q)(χγµγ

5χ)] interaction⇒ Cross section ∝ target spin

)2WIMP mass (GeV/c10 210

310 410

(p

b)

pSD

σ

410

310

210

110

1

cMSSM

DAMA/LIBRA 2008

COUPP 2007

COUPP 2011

ICECUBE 2011

KIMS 2007

SUPERKAMIOKANDE 2011SIM

PLE 2010

PICASSO 2009

PICASSO 2012

PICASSO arXiv:1202.1240

Note: CoGeNT & CRESST have very low sensitivity to spin-dependent DM scattering.Joachim Kopp Dark Matter and Dark Energy 36

Inelastic dark matter?

Idea: There may be two DM states χ and χ′ with

mχ′ = mχ + δ

Scattering proceeds via

χ+ N → χ′ + N

Modified kinematics compared to elastic scatteringAffects different target nuclei differently


Tucker-Smith Weiner hep-ph/0101138

Inelastic dark matter?

Idea: There may be two DM states χ and χ′ with

mχ′ = mχ + δ

30 40 5010-41

10-40

10-39

10-38

10-37

mΧ @GeVD

WIM

P-nu

cleo

ncr

oss

sect

ion

ΣSI

@cm2 D

iDM, ∆=90keV

Limits: 90%

Contours: 90%, 3Σ

v0 = 220 kms, vesc = 550 kms

Xenon

-100

KIM

SCRESST-II

DAMA

CR

ESST

comm

. run

plot from JK Schwetz Zupan 1110.2721


Tucker-Smith Weiner hep-ph/0101138

Isospin-violating dark matter?

Idea: Dark matter could couple differently to protons and neutrons⇒ Detection efficiencies of different target materials change

-1.2 -1 -0.8 -0.6fn / f

p

10-4

10-3

10-2

10-1

A2 ef

f / A

2

-1.2 -1 -0.8 -0.6

Si / Ca /

O

Na

Ge

I

Xe

W


fn, fp: DM couplings to protons and neutronsAeff: Effective nuclear mass for DM scattering


Feng Kumar Marfatia Sanford 1102.4331

Isospin-violating dark matter?

Idea: Dark matter could couple differently to protons and neutrons⇒ Detection efficiencies of different target materials change

101 102 103

10-41

10-40

10-39

10-38

10-37

mΧ @GeVD

WIM

P-nu

cleo

ncr

oss

sect

ion

Σef

f@cm

2 DIVDM, fn fp=-0.7

Limits: 90%

Contours: 90%, 3Σ

v0 = 220 kms, vesc = 550 kms

CDMS-SiCDMS low-thr

Xenon-100

KIMS

CRESST-II

DAMA

CRESST comm. run

CDMS



Feng Kumar Marfatia Sanford 1102.4331

Leptophilic dark matter?

Idea: DM could couple only to leptons at tree levelDAMA and CoGeNT do not reject electron-recoils as backgroundBut: Electron recoils above threshold (& 1 keV) strongly suppressed(electron needs large initial momentum→ probe high-p tail of wave functions)

105 106 10710-16

10-14

10-12

10-10

10-8

10-6

p @eVD

p2ÈΧHpL

2@e

V-

1 D

Iodine

1s

2p 2s3d

3p 3s4d

4p 4s

5p 5s

105 106 10710-16

10-14

10-12

10-10

10-8

10-6

p @eVD

p2ÈΧHpL

2@e

V-

1 D

Sodium

1s

2p

2s

3s

Thus: DM–nucleus scattering dominates, even if loop-inducedBut: Loop diagrams forbidden for some modelsProblems then:

I Very large couplings needed to compensate wave function suppressionI Poor fit to DAMA and CoGeNT energy spectra

0 5 10 15 20

-0.04

-0.02

0.00

0.02

0.04

0.06

E @keVD

Sm@d-1kg-1keV-1D

Full analysis

wo 1st bin

Σe0 = 5.´ 10-31 cm2, mΧ = 816. GeV

Χ2 = 55.9

Σe0 = 5.01´ 10-31 cm2, mΧ = 202. GeV

Χ2 = 20.6


Fox Poppitz arXiv:0811.0399JK Niro Schwetz Zupan arXiv:0907.3159


Idea: DM could couple only to leptons at tree levelDAMA and CoGeNT do not reject electron-recoils as backgroundBut: Electron recoils above threshold (& 1 keV) strongly suppressedThus: DM–nucleus scattering dominates, even if loop-induced

ℓ−

ℓ−

γ

χ

χ

p+

p+

q

γ

γ

N , pµ

N , p′µ

χ, kµ

χ, k′µ

ℓ−

ℓ−

q

γ

γ

N , pµ

N , p′µ

χ, kµ

χ, k′µ

ℓ−

ℓ−

But: Loop diagrams forbidden for some modelsProblems then:


0 5 10 15 20

-0.04

-0.02

0.00

0.02

0.04

0.06

E @keVD

Sm@d-1kg-1keV-1D

Full analysis

wo 1st bin

Σe0 = 5.´ 10-31 cm2, mΧ = 816. GeV

Χ2 = 55.9

Σe0 = 5.01´ 10-31 cm2, mΧ = 202. GeV

Χ2 = 20.6




Idea: DM could couple only to leptons at tree levelDAMA and CoGeNT do not reject electron-recoils as backgroundBut: Electron recoils above threshold (& 1 keV) strongly suppressedThus: DM–nucleus scattering dominates, even if loop-inducedBut: Loop diagrams forbidden for some modelse.g. axial vector couplings g2/M2(χγµγ5χ)(fγµγ5f )

Problems then:I Very large couplings needed to compensate wave function suppressionI Poor fit to DAMA and CoGeNT energy spectra

0 5 10 15 20

-0.04

-0.02

0.00

0.02

0.04

0.06

E @keVD

Sm@d-1kg-1keV-1D

Full analysis

wo 1st bin

Σe0 = 5.´ 10-31 cm2, mΧ = 816. GeV

Χ2 = 55.9

Σe0 = 5.01´ 10-31 cm2, mΧ = 202. GeV

Χ2 = 20.6




Idea: DM could couple only to leptons at tree levelDAMA and CoGeNT do not reject electron-recoils as backgroundBut: Electron recoils above threshold (& 1 keV) strongly suppressedThus: DM–nucleus scattering dominates, even if loop-inducedBut: Loop diagrams forbidden for some modelsProblems then:


0 5 10 15 20

-0.04

-0.02

0.00

0.02

0.04

0.06

E @keVD

Sm@d-1kg-1keV-1D

Full analysis

wo 1st bin

Σe0 = 5.´ 10-31 cm2, mΧ = 816. GeV

Χ2 = 55.9

Σe0 = 5.01´ 10-31 cm2, mΧ = 202. GeV

Χ2 = 20.6



Indirect DM detection — where to look

The Galactic Center

Pros:Highest DM density

Cons:DM distribution uncertainMany background sources

Dwarf Galaxies

Pros:Few backgrounds

Cons:Relatively low DM density


Indirect DM detection — where to look

The Galactic Center

Pros:Highest DM density

Cons:DM distribution uncertainMany background sources

Dwarf Galaxies

101 102 103

WIMP mass [GeV]

10-26

10-25

10-24

10-23

10-22

10-21

WIM

Pcr

oss

sect

ion

[cm3

/s]


3*10−26

bb Channel

τ+ τ− Channel

µ + µ− Channel

W+W− Channel




Cons:Relatively low DM density


Thermal relic cross-section

Fermi-LAT, 1108.3546see also Geringer-Sameth Koushiappas 1108.2914

Indirect DM detection — where to look (2)

Cosmic antimatter

Pros:Few background sources

Cons:Backgrounds uncertainPropagation of chargedparticle has large uncertaintiesNon-directional

High-energy neutrinos

Idea:DM capture/annihilation in the SunFlux dominated by capture rate


Cons:Low neutrino cross sections



Cosmic antimatter








PAMELA collaboration, 0810.4995


Cosmic antimatter








IceCube collaboration, 1111.2738PAMELA collaboration, 0810.4995

What is a sphaleron?

SU(2) gauge field vacuum configurations are classified according to theirwinding number (or Chern-Simons number)

NCS = 116π2

∫ t0dt

∫d3x tr Fµν Fµν

NCS =1

16π2

∫ t

0dt


Configurations with different winding number cannot be continuouslytransformed into each other.

Sphalerons are processes (with E > 0) that change the winding numberIn the SM, a change in winding number corresponds to a change inB + L. In fact, considering only left-handed (SU(2)L-charged) fermions:

jµB+L =∑ψ=q,`

12ψγµ(1− γ5)ψ

A change in B + L is equivalent to a change in NCS:



SU(2) gauge field vacuum configurations are classified according to theirwinding number (or Chern-Simons number) NCS = 1

16π2

∫ t0dt


Sphalerons are processes (with E > 0) that change the winding numberTheir energy is of order mH , the symmetry breaking scale (100 GeV)

In the SM, a change in winding number corresponds to a change inB + L. In fact, considering only left-handed (SU(2)L-charged) fermions:

jµB+L =∑ψ=q,`

12ψγµ(1− γ5)ψ




SU(2) gauge field vacuum configurations are classified according to theirwinding number (or Chern-Simons number) NCS = 1

16π2

∫ t0dt


Sphalerons are processes (with E > 0) that change the winding numberIn the SM, a change in winding number corresponds to a change inB + L. In fact, considering only left-handed (SU(2)L-charged) fermions:

jµB+L =∑ψ=q,`

12ψγµ(1− γ5)ψ


∂t

∫d3x j0B+L ≡

∫d3x

12∂µψγ

µγ5ψ (since ∂µψγµψ = 0)

= − 116π2

∫d3x tr Fµν Fµν (chiral anomaly)

= −∂tNCS


the missing 95%: theory and phenomenology of …home.fnal.gov/~jkopp/pdf/dmde.pdf(inferred from...

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