The standard model of cosmology: CDM - McGill patscott/talks/PS_Oslo.pdfThe standard model of...

Cosmological ModelsCDM – Background

CDM – Selected results

The standard model of cosmology: ΛCDM

Pat Scott

Department of Physics, McGill University

Slides available fromwww.physics.mcgill.ca/ patscott

Pat Scott – Oct 5 – University of Oslo The standard model of cosmology: ΛCDM

http://www.physics.mcgill.ca/~patscott



Outline

1 Cosmological Models2 CDM – Background

Evidence and modelsDark matter detection

3 CDM – Selected resultsIndirect detectionMultimessenger particle physics: global fitsEffects of dark matter on stars




The Friedmann Equation

Take the Einstein Equation from General Relativity:

Gµν = 8πTµν + Λgµν . (1)

Assume the Universe to be isotropic and homogeneous=⇒ Friedmann-Robertson Walker (FRW) metric:

gµνxµxν = dt2 + R(t)2(

dr2

1− kr2 + r2dΩ2). (2)

Solve µ = 0, ν = 0 of (1) =⇒ Friedmann Equation:

H(t) ≡ R(t)R(t)

=8πG

3ρ(t)− k

R(t)2 . (3)






H(t) ≡ R(t)R(t)

=8πG

3ρ(t)− k

R(t)2 . (4)

Solving this gives Hubble parameter H(t) for someenergy density-scalefactor relation R(t) = f (ρ(t))

curvature k ∈ +1,0,−1H(t) encodes the dynamic evolution of the Universe

Critical density:For a flat Universe k = 0. This defines

critical density: ρc = 3H(t)8πG

cosmological density: Ωx ≡ ρxρc






H(t) ≡ R(t)R(t)

=8πG

3ρ(t)− k

R(t)2 . (4)


curvature k ∈ +1,0,−1

H(t) encodes the dynamic evolution of the Universe









H(t) ≡ R(t)R(t)

=8πG

3ρ(t)− k

R(t)2 . (4)


curvature k ∈ +1,0,−1H(t) encodes the dynamic evolution of the Universe







Equations of state

Equations of state:1st law of thermodynamics (µ = 0 in conservation of Tµν) is

d(ρR3) = −pd(R3), i.e. ∆E = −p∆V (5)

with a constant equation of state ρ = wp, we get energydensity-scalefactor relations

ρ ∝ R−3(1+w) (6)

For different types of energy:Matter: w = 0 =⇒ ρ ∝ R−3

Radiation: w = 1/3 =⇒ ρ ∝ R−4

Vacuum (Λ): w = −1 =⇒ ρ ∝ constant

This is basically enough to solve the Friedmann Equation.Pat Scott – Oct 5 – University of Oslo The standard model of cosmology: ΛCDM



Ingredients of ΛCDM

Ingredients required Choices in ΛCDMfor a cosmological model

A theory of gravity GR+ associated assumptions +isotropy,homogeneity

Types of energy radiation,matter, vacuum/dark energy

their equations of state w = 1/3, 0, −1/other

their (self−)interactions photons,baryonic (SM) matter+cold dark matter (CDM), ??

An initial spectrum approximately scale invariantof perturbations on large scales




An aside: inflation

QuestionIsn’t inflation part of the ΛCDM model?

AnswerNot really, no.

Approximately scale-invariant spectrum of perturbations to startwith, on CMB scales (small k )? Yes.Due to inflation by definition? No.

Pδ(k) ∝ PR(k) ∝ kns−1

+α log k/k0

(7)

ΛCDM does not demand inflation, just as it does not demandany particular CDM−inflation is just an idea for getting the required spectrum on CMB scales−any particular DM model is just an idea for getting CDM




An aside: inflation

QuestionIsn’t inflation part of the ΛCDM model?

AnswerNot really, no.

Approximately scale-invariant spectrum of perturbations to startwith, on CMB scales (small k )? Yes.Due to inflation by definition? No.

Pδ(k) ∝ PR(k) ∝ kns−1+α log k/k0 (7)

ΛCDM does not demand inflation, just as it does not demandany particular CDM−inflation is just an idea for getting the required spectrum on CMB scales−any particular DM model is just an idea for getting CDM




Cosmological probes & ‘concordance cosmology’

Joint fit to multiplecosmological observablesgives a consistent set ofparameter values:

ΩΛ ≈ 0.73Ωmatter ≈ 0.27

=

ΩCDM ≈ 0.23 + Ωbaryons ≈ 0.04→ΛCDM

(I follow a similar global fitstrategy to hunt for DM andparticle theories beyond the SM)

23

Fit !M !k w

SNe 0.287+0.029+0.039!0.027!0.036 0 (fixed) -1 (fixed)

SNe+BAO 0.285+0.020+0.011!0.020!0.009 0 (fixed) !1.011+0.076+0.083

!0.082!0.087

SNe+CMB 0.265+0.022+0.018!0.021!0.016 0 (fixed) !0.955+0.060+0.059

!0.066!0.060

SNe+BAO+CMB 0.274+0.016+0.013!0.016!0.012 0 (fixed) !0.969+0.059+0.063

!0.063!0.066

SNe+BAO+CMB 0.285+0.020+0.011!0.019!0.011 !0.009+0.009+0.002

!0.010!0.003 -1 (fixed)

SNe+BAO+CMB 0.285+0.020+0.010!0.020!0.010 !0.010+0.010+0.006

!0.011!0.004 !1.001+0.069+0.080!0.073!0.082

TABLE 6Fit results on cosmological parameters !M, !k and w. The parameter values are followed by their statistical (!stat) andsystematic (!sys) uncertainties. The parameter values and their statistical errors were obtained from minimizing the "2 ofEq. 3. The fit to the SNe data alone results in a "2 of 310.8 for 303 degrees of freedom with a ""2 of less than one forthe other fits. The systematic errors were obtained from fitting with extra nuisance parameters according Eq. 5 and

subtracting from the resulting error, !w/sys, the statistical error: !sys = (!2w/sys

! !2stat)

1/2.

0.0 0.1 0.2 0.3 0.4 0.5

-1.5

-1.0

-0.5

0.0

mm

w 0.2 0.3 0.4

-1.0

-0.7

w

0.2 0.3 0.4

w/ sys

w/o NB99

-1.3

-1.0

-0.7

w

-1.3

SNe

BAO

CMB

Fig. 14.— 68.3 %, 95.4 % and 99.7% confidence level contours onw and !M, for a flat Universe. The top plot shows the individualconstraints from CMB, BAO and the Union SN set, as well as thecombined constraints (filled gray contours, statistical errors only).The upper right plot shows the e#ect of including systematic errors.The lower right plot illustrates the impact of the SCP Nearby 1999data.

straints from combining SNe, CMB and BAO are consis-tent with a flat !CDM Universe (as seen in Table 6). Fig.15 shows the corresponding constraints in the "M ! "!

plane.Finally, one can attempt to investigate constraints on

a redshift dependent equation of state (EOS) parameterw(z). Initially we consider this in terms of

w(z) = w0 + waz

1 + z, (10)

shown by Linder (2003) to provide excellent approxima-tion to a wide variety of scalar field and other dark en-ergy models. Later, we examine other aspects of timevariation of the dark energy EOS. Assuming a flat Uni-verse and combining the Union set with constraints fromCMB, we obtain constraints on w0, the present valueof the EOS, and wa, giving a measure of its time vari-ation, as shown in Fig. 16. (A cosmological constanthas w0 = !1, wa = 0.) Due to degeneracies within theEOS and between the EOS and the matter density "M,the SN dataset alone does not give appreciable leverageon the dark energy properties. By adding other mea-surements, the degeneracies can be broken and currentlymodest cosmology constraints obtained.

Fig. 16 (left) shows the combination of the SN datawith either the CMB constraints or the BAO constraints.

0.0 0.5 1.0

0.0

0.5

1.0

1.5

2.0

FlatBAO

CMB

SNe

No Big Bang

Fig. 15.— 68.3 %, 95.4 % and 99.7% confidence level contourson !! and !M obtained from CMB, BAO and the Union SN set,as well as their combination (assuming w = !1).

The results are similar; note that including either one re-sults in a sharp cut-o# at w0+wa = 0, from the physics asmentioned in regards to Eq. 9. Since w(z " 1) = w0+wa

in this parameterization, any model with more positivehigh-redshift w will not yield a matter-dominated earlyUniverse, altering the sound horizon in conflict with ob-servations.

Note that BAO do not provide a purely “low” redshiftconstraint, because implicit within the BAO data anal-ysis, and hence the constraint, is that the high redshiftUniverse was matter dominated (so the sound horizon

Kowalski et al ApJ 2008





Outline








How we know dark matter exists

The only way to consistently explain:1 rotation curves + vel. dispersions2 gravitational lensing3 cosmological data

Large-scale structure (2dF/Chandra/SDSS-BAO) saysΩmatter ≈ 0.27BBN says that Ωbaryonic ≈ 0.04=⇒ Ωnon−baryonic ≈ 5× ΩbaryonsCMB (WMAP) and SN1a agree; also indicate that Ωtotal ≈ 1=⇒ universe is 23% dark matter, 4% baryonic (visible)matter, 73% something else


(Clowe et al., ApJL 2006)




What we know about it

Must be:massive (gravitationally-interacting)

unable to interact via the electromagnetic force (dark)

non-baryonic

“cold(ish)” (in order to allow structure formation)

stable on cosmological timescales

produced with the right relic abundance in the early Universe.

Good options:Weakly Interacting Massive Particles (WIMPs)

sterile neutrinos

gravitinos

axions

axinos

hidden sector dark matter (e.g. WIMPless dark matter)Pat Scott – Oct 5 – University of Oslo The standard model of cosmology: ΛCDM




What we know about it

Must be:massive (gravitationally-interacting)

unable to interact via the electromagnetic force (dark)

non-baryonic

“cold(ish)” (in order to allow structure formation)

stable on cosmological timescales

produced with the right relic abundance in the early Universe.

Good options:Weakly Interacting Massive Particles (WIMPs)

sterile neutrinos

gravitinos

axions

axinos

hidden sector dark matter (e.g. WIMPless dark matter)Pat Scott – Oct 5 – University of Oslo The standard model of cosmology: ΛCDM

Bad options:primordial black holes

MAssive Compact Halo Objects (MACHOs)

standard model neutrinos




WIMPs at a glance

Dark because no electromagnetic interactionsCold because very massive (∼10 GeV to ∼10 TeV)Non-baryonic and stable - no problems with BBN or CMBWeak-scale annihilation cross-sections naturally lead to arelic abundance of the right order of magnitude


(Kolb & Turner1990)




WIMPs at a glance

Many theoretically well-motivated particle candidatesSupersymmetric (SUSY) neutralinos χ if R-parity is conserved -lightest mixture of neutral higgsinos and gauginosInert Higgses - extra Higgs in the Standard ModelKaluza-Klein particles - extra dimensionsright-handed neutrinos, sneutrinos, other exotic things. . .

Weak interaction means scattering with nuclei→ detectionchannelMany WIMPs are Majorana particles (own antiparticles)=⇒ self-annihilation cross-section





Outline








Ways to detect WIMPs

Direct detection – nuclear collisions and recoils – CDMS,XENON, DAMA, CRESST, CoGeNT

Direct production – missing ET or otherwise – LHC,TevatronIndirect detection – annihilations producing

gamma-rays – Fermi, HESS, CTAanti-protons – PAMELA, AMSanti-deuterons – GAPSneutrinos – IceCube, ANTARESe+e− – PAMELA, Fermi, ATIC, AMS→ secondary radiation: Compton−1,synchrotron, bremsstrahlungsecondary impacts on the CMB

Dark stars – JWST, VLT






Direct detection – nuclear collisions and recoils – CDMS,XENON, DAMA, CRESST, CoGeNTDirect production – missing ET or otherwise – LHC,Tevatron

Indirect detection – annihilations producinggamma-rays – Fermi, HESS, CTAanti-protons – PAMELA, AMSanti-deuterons – GAPSneutrinos – IceCube, ANTARESe+e− – PAMELA, Fermi, ATIC, AMS→ secondary radiation: Compton−1,synchrotron, bremsstrahlungsecondary impacts on the CMB







Direct detection – nuclear collisions and recoils – CDMS,XENON, DAMA, CRESST, CoGeNTDirect production – missing ET or otherwise – LHC,TevatronIndirect detection – annihilations producing

gamma-rays – Fermi, HESS, CTAanti-protons – PAMELA, AMSanti-deuterons – GAPSneutrinos – IceCube, ANTARESe+e− – PAMELA, Fermi, ATIC, AMS→ secondary radiation: Compton−1,synchrotron, bremsstrahlungsecondary impacts on the CMB





Indirect detectionMultimessenger particle physics: global fitsEffects of dark matter on stars

Outline








Finding dark matter with neutrino telescopes

The cartoon version:

1 Halo WIMPs crash into the Sun2 Some lose enough energy in the scatter to

be gravitationally bound3 Scatter some more, sink to the core4 Annihilate with each other, producing

neutrinos5 Propagate+oscillate their way to the Earth,

convert into muons in ice/water6 Look for Cerenkov radiation from the

muons in IceCube, ANTARES, etc






The cartoon version:1 Halo WIMPs crash into the Sun

2 Some lose enough energy in the scatter tobe gravitationally bound

3 Scatter some more, sink to the core4 Annihilate with each other, producing









The cartoon version:1 Halo WIMPs crash into the Sun2 Some lose enough energy in the scatter to

be gravitationally bound

3 Scatter some more, sink to the core4 Annihilate with each other, producing










be gravitationally bound3 Scatter some more, sink to the core

4 Annihilate with each other, producingneutrinos

5 Propagate+oscillate their way to the Earth,convert into muons in ice/water

6 Look for Cerenkov radiation from themuons in IceCube, ANTARES, etc








neutrinos











convert into muons in ice/water






The IceCube Neutrino Observatory

86 strings1.5–2.5 km deep inAntarctic ice sheet∼125 m spacingbetween strings∼70 m in DeepCore(10× higher opticaldetector density)1 km3 instrumentedvolume (1 Gton)





The IceCube Neutrino Observatory

86 strings1.5–2.5 km deep inAntarctic ice sheet∼125 m spacingbetween strings∼70 m in DeepCore(10× higher opticaldetector density)1 km3 instrumentedvolume (1 Gton)


MeanTrue value

Scott, Danninger, Savage, Edsjo & Hultqvist 2012

0.2

0.4

0.6

0.8

1.0

0 200 400 600 800mχ0

1(GeV)

20

40

60

80

100

120

140

Pre

dict

edsi

gnal

even

ts

Relative

probabilityP/P

max

IC22× 100flat priorsCMSSM µ > 0Marg. posterior

Mean

Scott, Danninger, Savage, Edsjo & Hultqvist 2012

0.2

0.4

0.6

0.8

1.0

0 200 400 600 800mχ0

1(GeV)

−43

−42

−41

−40

−39

−38lo

g 10

( σSD,p/c

m2)

Relative

probabilityP/P

max

IC22× 100flat priorsCMSSM µ > 0Marg. posterior

PS, Savage, Edsjö & TheIceCube Collaboration,arxiv:1207.0810




Gamma-rays from dark matter

2 photons (or Z+photon):

monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z

Internal bremsstrahlung:

hard gamma-ray spectrum

Secondary decay:

soft(er) continuum spectrum

χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ


monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ


monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ

3 main gamma-ray channels:

monochromatic linesinternal bremsstrahlung (FSR + VIB)continuum from secondary decay







monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ


monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ


monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ

3 main gamma-ray channels:monochromatic lines

internal bremsstrahlung (FSR + VIB)continuum from secondary decay







monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ


monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ


monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ

3 main gamma-ray channels:monochromatic linesinternal bremsstrahlung (FSR + VIB)

continuum from secondary decay







monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ


monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ


monochromatic lines

χ 0

1

χ 0

1

γ

γ/Z



Secondary decay:


χ 0

1

χ 0

1

γ

χ 0

1

χ 0

1

SM

SM

SM

SM

SM

SM

SM

SM

π

γ

γ

γ

3 main gamma-ray channels:monochromatic linesinternal bremsstrahlung (FSR + VIB)continuum from secondary decay





Targets

Φ ∝ annihilation rate ∝ ρ2DM

Likely targets:Galactic centre - large signal, large BGGalactic halo - moderate signal, moderate BGdwarf galaxies - low statistics, low BGclusters/extragalactic diffuse - large modellinguncertainties, low signal, low BGdark clumps - low statistics, low BG





Targets


Likely targets:Galactic centre - large signal, large BG

Galactic halo - moderate signal, moderate BGdwarf galaxies - low statistics, low BGclusters/extragalactic diffuse - large modellinguncertainties, low signal, low BGdark clumps - low statistics, low BG


0

5

10

15

20

25

30

35

40

Counts

p-value=0.85, χ 2red=14.3/21

Signal counts: 53.4 (4.26σ) 80.5 - 208.5 GeV

Reg3 (ULTRACLEAN), Eγ =129.6 GeV

100 150 200

E [GeV]

-10

0

10

Counts - Model

Weniger, arXiv:1204.1971Hooper & Linden, arXiv:1110.0006




Targets


Likely targets:Galactic centre - large signal, large BGGalactic halo - moderate signal, moderate BGdwarf galaxies - low statistics, low BG

clusters/extragalactic diffuse - large modellinguncertainties, low signal, low BGdark clumps - low statistics, low BG

Pat Scott – Oct 5 – University of Oslo The standard model of cosmology: ΛCDM101 102 103

WIMP mass [GeV]

10-26

10-25

10-24

10-23

10-22

10-21

10-20

10-19

WIM

P c

ross

sect

ion [

cm3

/s]

Upper limits, bb channel

3 ·10−26

Bootes I

Carina

Coma Berenices

Draco

Fornax

Sculptor

Segue 1

Sextans

Ursa Major II

Ursa Minor

Joint Likelihood, 10 dSphs

mχ1

0 (TeV)

log 10

[ <

σ v>

(cm

3 s−

1 ) ]

Profile likelihood

Flat priors

CMSSM, µ>0

Segue 9 mth + All, BF=50

Scott et al. 2009

0 0.5 1−30

−28

−26

−24

Fermi-LAT Collab., Phys. Rev. Lett 2011PS, Conrad, Edsjö et al, JCAP 2010




Targets


Likely targets:Galactic centre - large signal, large BGGalactic halo - moderate signal, moderate BGdwarf galaxies - low statistics, low BGclusters/extragalactic diffuse - large modellinguncertainties, low signal, low BGdark clumps - low statistics, low BG





An example of dark clumps: Ultracompact minihalos

QuestionWhat is an ultracompact minihalo (UCMH)?

AnswerA DM halo that collapses shortly after matter-radiation equality

‘Shortly’ means zcollapse is O(100) or more (vs zeq ∼ 3000)

=⇒ isolated collapse

=⇒ formation by radial infall

=⇒ very steep density profile→ ρ ∝ r−9/4

=⇒ excellent indirect detection targets

Also good lensing prospects Ricotti & Gould ApJ 2009; Li et al Phys. Rev. D 2012


Scott & SivertssonPhys. Rev. Lett. 2009Lacki & Beacom ApJL 2010




An example of dark clumps: Ultracompact minihalos

QuestionHow would UCMHs be created?

AnswerLarge amplitude density perturbations in the early Universe(e.g. on small scales)

Small-scale power in primordial perturbation spectrum(e.g. features in the potential associated with inflation)

Phase transitions

Other seeds (e.g. cosmic strings)





Primordial power spectrum

Limits on PR from gamma-ray searches for UCMHs ∼5 ordersbetter than from PBHs=⇒ strong limits on inflationary models

k (Mpc−1)

P δ(k

)

WIMP kinetic decoupling

P R(k

)

10−9

10−8

10−7

10−6

10−5

10−4

10−3

10−2

10−1

10−3

10−2

10−1 1 10 102 103 104 105 106 107 108 109 101

0101

1101

2101

3101

4101

5101

6101

7101

8101

9

10−9

10−8

10−7

10−6

10−5

10−4

10−3

10−2

Allowed regions

Ultracompact minihalos (gamma rays, Fermi -LAT)

Ultracompact minihalos (reionisation, WMAP5 τe)

Primordial black holes

CMB, Lyman-α, LSS and other cosmological probes


Bringmann, PS, Akrami, Phys. Rev. D 2012




Implications for cosmology

Impacts on inflation:

Notvisible

Visiblein γ rays

Slow rollWMAP5

Slow rollWMAP7+SPT

Scale-free spectrumWMAP7 best fit

zc = 50

zc = 200

Scott, Adams, Bringmann & Easther 2012

0.85 0.90 0.95 1.00 1.05 1.10ns

−0.04

−0.03

−0.02

−0.01

0.00

0.01

0.02

0.03

α≡

[ dn

sdlnk

]

UCMHs

mχ = 10 GeV

mχ = 1 TeV

Limits on non-Gaussianities:

k = 2× 103 Mpc−1

Hierarchical Scaling

Excluded by UCMHszc = 200 zc = 1000

Excludedby PBHs

App

roxi

mat

esc

ale-

inva

rian

tam

plit

ude

−8 −6 −4 −2log10 [Pδ,Gaussian (k)]

−0.2

0.0

0.2

M3(k

)


PS, Adams, Bringmann, Easther, in prep.Shandera, Erickcek, PS, in prep




Outline








Putting it all together

Base Observables

+ XENON-100

(relic density, B-physics, LEP, etc.)

Grey contours correspond to Base Observables only

Mean

Scott, Danninger, Savage, Edsjo, Hultqvist & The IceCube Collab. (2012)

0.2

0.4

0.6

0.8

1.0

0.0 0.5 1.0 1.5 2.0 2.5m 1

2(TeV)

1

2

3

m0

(TeV

)

Relative

probabilityP/P

maxOld data

flat priorsCMSSM µ > 0Marg. posterior

Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.2 0.4 0.6 0.8 1.0mχ0

1(TeV)

−46

−45

−44

−43

−42

log 1

0

( σSI,p/c

m2)

Relative

probabilityP/P

max

Old dataflat priorsCMSSM µ > 0Marg. posterior

Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.2 0.4 0.6 0.8 1.0mχ0

1(TeV)

−43

−42

−41

−40

−39

−38

log 1

0

( σSD,p/c

m2)

Relative

probabilityP/P

max

Old dataflat priorsCMSSM µ > 0Marg. posterior

Contours indicate 1σ and 2σ credible regionsShading+contours indicate relative probability only, not overall goodness of fit

IceCube-22 with 100× boosted effective area (kinda like final IceCube)


PS, Savage, Edsjö & The IceCubeCollaboration, arxiv:1207.0810





Base Observables + XENON-100

+ CMS 5 fb−1+ IC22×100


Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.5 1.0 1.5 2.0 2.5m 1

2(TeV)

1

2

3

m0

(TeV

)

Relative

probabilityP/P

maxXENON100


Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.2 0.4 0.6 0.8 1.0mχ0

1(TeV)

−46

−45

−44

−43

−42

log 1

0

( σSI,p/c

m2)

Relative

probabilityP/P

max

XENON100flat priorsCMSSM µ > 0Marg. posterior

Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.2 0.4 0.6 0.8 1.0mχ0

1(TeV)

−43

−42

−41

−40

−39

−38

log 1

0

( σSD,p/c

m2)

Relative

probabilityP/P

max

XENON100flat priorsCMSSM µ > 0Marg. posterior









Base Observables + XENON-100+ CMS 5 fb−1

+ IC22×100


Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.5 1.0 1.5 2.0 2.5m 1

2(TeV)

1

2

3

m0

(TeV

)

Relative

probabilityP/P

max

XENON100+CMS 5 fb−1


Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.2 0.4 0.6 0.8 1.0mχ0

1(TeV)

−46

−45

−44

−43

−42

log 1

0

( σSI,p/c

m2)

Relative

probabilityP/P

max



Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.2 0.4 0.6 0.8 1.0mχ0

1(TeV)

−43

−42

−41

−40

−39

−38

log 1

0

( σSD,p/c

m2)

Relative

probabilityP/P

max











Base Observables + XENON-100+ CMS 5 fb−1+ IC22×100Grey contours correspond to Base Observables only

Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.5 1.0 1.5 2.0 2.5m 1

2(TeV)

1

2

3

m0

(TeV

)

Relative

probabilityP/P

max

IC22× 100+XENON100+CMS 5 fb−1


Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.2 0.4 0.6 0.8 1.0mχ0

1(TeV)

−46

−45

−44

−43

−42

log 1

0

( σSI,p/c

m2)

Relative

probabilityP/P

max



Mean


0.2

0.4

0.6

0.8

1.0

0.0 0.2 0.4 0.6 0.8 1.0mχ0

1(TeV)

−43

−42

−41

−40

−39

−38

log 1

0

( σSD,p/c

m2)

Relative

probabilityP/P

max










Outline








Reminder:










Reminder:

The cartoon version:1 Halo WIMPs crash into the Sun stars2 Some lose enough energy in the scatter to


neutrinos







Reminder:



neutrinos + other energetic particles







Reminder:



neutrinos + other energetic particles5 Particles dump their energy in the stellar

core






Reminder:



neutrinos + other energetic particles5 Particles dump their energy in the stellar

core6 Stellar structure responds, star evolves

accordingly





Stellar evolution with dark matter annihilation

Effective (surface) temperature, log10

(TeffK

)

Lum

inos

ity,

log 1

0(L

/L

)

Z = 0.01

3.653.703.75

−0.

6−

0.4

−0.

20.

00.

20.

40.

6

log10

( ρχ

GeV cm−3

)= −5

log10

( ρχ

GeV cm−3

)= 9

log10

( ρχ

GeV cm−3

)= 10

ZAMS





Finding ‘dark stars’

Best candidates have low masses, near Galactic CentreFirst stars also good targetsMaybe visible with JWST (if lensed), or by impacts onreionisationDARKSTARS stellar evolution code publicly available fromhttp://www.physics.mcgill.ca/ patscott/darkstars

10−1

100

101

1030

1031

1032

1033

1034

1035

Wavelength (µm)

Fλ (

W µ

m−

1 )

Planck 1σ, 14 months

WMAP 1σ, 7 years

0.2

0.4

0.6

0.8

1.0

Dar

kst

arfr

acti

onf D

S

0 100 200 300 400 500

DSP lifetime tDSP (Myr)

PS, Venkatesan, Roebber et al ApJ 2011Zackrisson, PS, Rydberg et al ApJ 2010; MNRAS Lett. 2010PS, Fairbairn, Edsjö, MNRAS 2009Fairbairn, PS, Edsjö, Phys. Rev. D 2008


http://www.physics.mcgill.ca/~patscott/darkstars




Summary

ΛCDM currently rests on CDM being some new particle –but what??There are many complementary ways to find out!Indirect detection generally probes masses andannihilation channelsStellar evolution can test both annihilation and interactionswith quarksUltracompact minihalos present an exciting way to alsoprobe cosmology at the same timeThe different probes can (and should) be put together intoglobal fits to gain a consistent picture. This will berequired for a credible detection to be claimed!





Extras 1: DarkStars code

Lots of options and switches: different velocitydistributions, widths, stellar orbits, WIMP conductivetransport / internal distribution schemes, particle data,stellar masses and metallicities, numerical options. . .Save and restart - good for evolving part-way then tryingdifferent late-stage scenariosDARKSTARS 2.0 coming soon: conversion to full Z = 0(new opacities, equation of state) – DARKSTARS 1.03 canonly do Z = 0 on pre-MSFuture options for expansion to include alternative formfactors and/or WIMP evaporationDARKSTARS 1.03 publicly available fromhttp://www.physics.mcgill.ca/ patscott/darkstars


http://www.physics.mcgill.ca/~patscott/darkstars




Extras 2: CMSSM

Model: focus has mostly been on the Constrained MSSM(CMSSM)

GUT boundary conditions on soft SUSY breaking parameters such thatonly 4 free parameters and 1 sign remain

includes the simplest implementation of mSUGRA

m0 scalar mass parameterm 1

2gaugino mass parameter

tan β ratio of Higgs VEVsA0 trilinear couplingsgn µ Higgs mass parameter

(+ve in our scans)

Just a testbed framework – all techniques are applicable to anyMSSM parameterisation


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