making dark matter - uni-bonn.de · making dark matter manuel drees bonn university & bethe...
TRANSCRIPT
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Making Dark MatterManuel Drees
Bonn University & Bethe Center for Theoretical Physics
Making Dark Matter – p. 1/35
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Contents
1 Introduction: The need for DM
Making Dark Matter – p. 2/35
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Contents
1 Introduction: The need for DM
2 Particle Dark Matter candidates
Making Dark Matter – p. 2/35
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Contents
1 Introduction: The need for DM
2 Particle Dark Matter candidates
3 Making particle DM
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Contents
1 Introduction: The need for DM
2 Particle Dark Matter candidates
3 Making particle DMa) WIMPs at low TR
Making Dark Matter – p. 2/35
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Contents
1 Introduction: The need for DM
2 Particle Dark Matter candidates
3 Making particle DMa) WIMPs at low TR
b) Thermal WIMPs (freeze out)
Making Dark Matter – p. 2/35
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Contents
1 Introduction: The need for DM
2 Particle Dark Matter candidates
3 Making particle DMa) WIMPs at low TR
b) Thermal WIMPs (freeze out)
c) FIMPs (freeze in)
Making Dark Matter – p. 2/35
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Contents
1 Introduction: The need for DM
2 Particle Dark Matter candidates
3 Making particle DMa) WIMPs at low TR
b) Thermal WIMPs (freeze out)
c) FIMPs (freeze in)
d) Thermal gravitino production
Making Dark Matter – p. 2/35
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Contents
1 Introduction: The need for DM
2 Particle Dark Matter candidates
3 Making particle DMa) WIMPs at low TR
b) Thermal WIMPs (freeze out)
c) FIMPs (freeze in)
d) Thermal gravitino production
e) Production in inflaton decay
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Contents
1 Introduction: The need for DM
2 Particle Dark Matter candidates
3 Making particle DMa) WIMPs at low TR
b) Thermal WIMPs (freeze out)
c) FIMPs (freeze in)
d) Thermal gravitino production
e) Production in inflaton decay
4 Summary
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Introduction: the need for Dark Matter
Several observations indicate existence of non-luminousDark Matter (DM) (more exactly: missing force)
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Introduction: the need for Dark Matter
Several observations indicate existence of non-luminousDark Matter (DM) (more exactly: missing force)
Galactic rotation curves imply ΩDMh2 ≥ 0.05.
Ω: Mass density in units of critical density; Ω = 1 means flatUniverse.h: Scaled Hubble constant. Observation: h = 0.72 ± 0.07
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Introduction: the need for Dark Matter
Several observations indicate existence of non-luminousDark Matter (DM) (more exactly: missing force)
Galactic rotation curves imply ΩDMh2 ≥ 0.05.
Ω: Mass density in units of critical density; Ω = 1 means flatUniverse.h: Scaled Hubble constant. Observation: h = 0.72 ± 0.07
Models of structure formation, X ray temperature ofclusters of galaxies, . . .
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Introduction: the need for Dark Matter
Several observations indicate existence of non-luminousDark Matter (DM) (more exactly: missing force)
Galactic rotation curves imply ΩDMh2 ≥ 0.05.
Ω: Mass density in units of critical density; Ω = 1 means flatUniverse.h: Scaled Hubble constant. Observation: h = 0.72 ± 0.07
Models of structure formation, X ray temperature ofclusters of galaxies, . . .
Cosmic Microwave Background anisotropies (WMAPetc.) imply ΩDMh2 = 0.112 ± 0.006 PDG, 2012 edition
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Need for non–baryonic DM
Total baryon density is determined by:
Big Bang Nucleosynthesis
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Need for non–baryonic DM
Total baryon density is determined by:
Big Bang Nucleosynthesis
Analyses of CMB data
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Need for non–baryonic DM
Total baryon density is determined by:
Big Bang Nucleosynthesis
Analyses of CMB data
Consistent result: Ωbarh2 ≃ 0.02
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Need for non–baryonic DM
Total baryon density is determined by:
Big Bang Nucleosynthesis
Analyses of CMB data
Consistent result: Ωbarh2 ≃ 0.02
=⇒ Need non–baryonic DM!
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Need for exotic particles
Only possible non–baryonic particle DM in SM: Neutrinos!
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Need for exotic particles
Only possible non–baryonic particle DM in SM: Neutrinos!
Make hot DM: do not describe structure formation correctly=⇒ Ωνh
2 ≤ 0.0062
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Need for exotic particles
Only possible non–baryonic particle DM in SM: Neutrinos!
Make hot DM: do not describe structure formation correctly=⇒ Ωνh
2 ≤ 0.0062
=⇒ Need exotic particles as DM!
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Need for exotic particles
Only possible non–baryonic particle DM in SM: Neutrinos!
Make hot DM: do not describe structure formation correctly=⇒ Ωνh
2 ≤ 0.0062
=⇒ Need exotic particles as DM!
Possible loophole: primordial black holes; not easy to makein sufficient quantity sufficiently early.
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What we need
Since h2 ≃ 0.5: Need ∼ 20% of critical density in
Matter (with negligible pressure, w ≃ 0)
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What we need
Since h2 ≃ 0.5: Need ∼ 20% of critical density in
Matter (with negligible pressure, w ≃ 0)
which still survives today (lifetime τ ≫ 1010 yrs)
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What we need
Since h2 ≃ 0.5: Need ∼ 20% of critical density in
Matter (with negligible pressure, w ≃ 0)
which still survives today (lifetime τ ≫ 1010 yrs)
and does not couple to elm radiation
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Remarks
Precise “WMAP” determination of DM density hingeson assumption of “standard cosmology”, includingassumption of nearly scale–invariant primordialspectrum of density perturbations: almost assumesinflation!
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Remarks
Precise “WMAP” determination of DM density hingeson assumption of “standard cosmology”, includingassumption of nearly scale–invariant primordialspectrum of density perturbations: almost assumesinflation!
Evidence for ΩDM >∼ 0.2 much more robust than that!(Does, however, assume standard law of gravitation.)
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Particle DM: Possibilities
Theorist’s tasks:
Introduce right kind of particle (stable, neutral,non–relativistic)
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Particle DM: Possibilities
Theorist’s tasks:
Introduce right kind of particle (stable, neutral,non–relativistic)
Make enough (but not too much) of it in early universe
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Particle DM: Possibilities
Theorist’s tasks:
Introduce right kind of particle (stable, neutral,non–relativistic)
Make enough (but not too much) of it in early universe
There are many possible ways to solve these tasks!
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Particle DM: Possibilities
Theorist’s tasks:
Introduce right kind of particle (stable, neutral,non–relativistic)
Make enough (but not too much) of it in early universe
There are many possible ways to solve these tasks!
=⇒ Use theoretical “prejudice” as guideline: Only considercandidates that solve (at least) one additional problem!
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Particle Candidates 1
Axion a: Pseudoscalar pseudo-Goldstone boson
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Particle Candidates 1
Axion a: Pseudoscalar pseudo-Goldstone bosonIntroduced to solve strong CP problem
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Particle Candidates 1
Axion a: Pseudoscalar pseudo-Goldstone bosonIntroduced to solve strong CP problem
Mass ma <∼ 10−3 eV
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Particle Candidates 1
Axion a: Pseudoscalar pseudo-Goldstone bosonIntroduced to solve strong CP problem
Mass ma <∼ 10−3 eV
Direct detection difficult, but possible
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Particle Candidates 1
Axion a: Pseudoscalar pseudo-Goldstone bosonIntroduced to solve strong CP problem
Mass ma <∼ 10−3 eV
Direct detection difficult, but possible
Neutralino χ : Majorana spin–1/2 fermion
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Particle Candidates 1
Axion a: Pseudoscalar pseudo-Goldstone bosonIntroduced to solve strong CP problem
Mass ma <∼ 10−3 eV
Direct detection difficult, but possible
Neutralino χ : Majorana spin–1/2 fermionRequired in supersymmetrized SM
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Particle Candidates 1
Axion a: Pseudoscalar pseudo-Goldstone bosonIntroduced to solve strong CP problem
Mass ma <∼ 10−3 eV
Direct detection difficult, but possible
Neutralino χ : Majorana spin–1/2 fermionRequired in supersymmetrized SM50 GeV <∼ mχ <∼ 1 TeV
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Particle Candidates 1
Axion a: Pseudoscalar pseudo-Goldstone bosonIntroduced to solve strong CP problem
Mass ma <∼ 10−3 eV
Direct detection difficult, but possible
Neutralino χ : Majorana spin–1/2 fermionRequired in supersymmetrized SM50 GeV <∼ mχ <∼ 1 TeV
Direct detection probably difficult, but possible
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Particle Candidates 2
Gravitino G: Majorana spin–3/2 fermion
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Particle Candidates 2
Gravitino G: Majorana spin–3/2 fermionRequired in supersymmetrized theory of gravity
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Particle Candidates 2
Gravitino G: Majorana spin–3/2 fermionRequired in supersymmetrized theory of gravity100 eV <∼ mG
<∼ 1 TeV
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Particle Candidates 2
Gravitino G: Majorana spin–3/2 fermionRequired in supersymmetrized theory of gravity100 eV <∼ mG
<∼ 1 TeV
Direct detection is virtually impossible
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Particle Candidates 2
Gravitino G: Majorana spin–3/2 fermionRequired in supersymmetrized theory of gravity100 eV <∼ mG
<∼ 1 TeV
Direct detection is virtually impossible
Here: focus on WIMP (e.g. Neutralino) and Gravitino.
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Making DM
Principal possibilities:
DM was in thermal equilibrium:
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Making DM
Principal possibilities:
DM was in thermal equilibrium:Implies lower bounds on temperature TR and χproduction cross section
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Making DM
Principal possibilities:
DM was in thermal equilibrium:Implies lower bounds on temperature TR and χproduction cross sectionΩ depends on particle physics and expansion history[Hubble parameter H(T ) ]
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Making DM
Principal possibilities:
DM was in thermal equilibrium:Implies lower bounds on temperature TR and χproduction cross sectionΩ depends on particle physics and expansion history[Hubble parameter H(T ) ]Example: Neutralino χ
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Making DM
Principal possibilities:
DM was in thermal equilibrium:Implies lower bounds on temperature TR and χproduction cross sectionΩ depends on particle physics and expansion history[Hubble parameter H(T ) ]Example: Neutralino χ
DM production from thermal plasma:
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Making DM
Principal possibilities:
DM was in thermal equilibrium:Implies lower bounds on temperature TR and χproduction cross sectionΩ depends on particle physics and expansion history[Hubble parameter H(T ) ]Example: Neutralino χ
DM production from thermal plasma:
Thermal equilibrium may never have been achieved(low TR and/or low interaction rate)
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Making DM
Principal possibilities:
DM was in thermal equilibrium:Implies lower bounds on temperature TR and χproduction cross sectionΩ depends on particle physics and expansion history[Hubble parameter H(T ) ]Example: Neutralino χ
DM production from thermal plasma:
Thermal equilibrium may never have been achieved(low TR and/or low interaction rate)Ω depends on particle physics, TR and H(T )
Making Dark Matter – p. 11/35
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Making DM
Principal possibilities:
DM was in thermal equilibrium:Implies lower bounds on temperature TR and χproduction cross sectionΩ depends on particle physics and expansion history[Hubble parameter H(T ) ]Example: Neutralino χ
DM production from thermal plasma:
Thermal equilibrium may never have been achieved(low TR and/or low interaction rate)Ω depends on particle physics, TR and H(T )
Example: Gravitino G with mG > 0.1 keV
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Making DM (cont.’d)
Non–thermal production:
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Making DM (cont.’d)
Non–thermal production:
From decay of heavier particle (χ, G)
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Making DM (cont.’d)
Non–thermal production:
From decay of heavier particle (χ, G)During phase transition (axion a)
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Making DM (cont.’d)
Non–thermal production:
From decay of heavier particle (χ, G)During phase transition (axion a)
During (p)reheating at end of inflation (G, χ)
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Making DM (cont.’d)
Non–thermal production:
From decay of heavier particle (χ, G)During phase transition (axion a)
During (p)reheating at end of inflation (G, χ)Depends strongly on details of particle physics andcosmology
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Making DM (cont.’d)
Non–thermal production:
From decay of heavier particle (χ, G)During phase transition (axion a)
During (p)reheating at end of inflation (G, χ)Depends strongly on details of particle physics andcosmology
Via particle–antiparticle asymmetry:
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Making DM (cont.’d)
Non–thermal production:
From decay of heavier particle (χ, G)During phase transition (axion a)
During (p)reheating at end of inflation (G, χ)Depends strongly on details of particle physics andcosmology
Via particle–antiparticle asymmetry:Assume symmetric contribution annihilates away:only “particles” left (see: baryons)
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Making DM (cont.’d)
Non–thermal production:
From decay of heavier particle (χ, G)During phase transition (axion a)
During (p)reheating at end of inflation (G, χ)Depends strongly on details of particle physics andcosmology
Via particle–antiparticle asymmetry:Assume symmetric contribution annihilates away:only “particles” left (see: baryons)If same mechanism generates baryon asymmetry:“Naturally” explains ΩDM ≃ 5Ωbaryon, if mχ ≃ 5mp
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Making DM (cont.’d)
Non–thermal production:
From decay of heavier particle (χ, G)During phase transition (axion a)
During (p)reheating at end of inflation (G, χ)Depends strongly on details of particle physics andcosmology
Via particle–antiparticle asymmetry:Assume symmetric contribution annihilates away:only “particles” left (see: baryons)If same mechanism generates baryon asymmetry:“Naturally” explains ΩDM ≃ 5Ωbaryon, if mχ ≃ 5mp
For WIMPs: Order of magnitude of ΩDM isunderstood; Ωbaryon isn’t
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Thermal history of the Universe
Currently: Universe dominated by dark energy (∼ 75%)and non–relativistic matter (∼ 25%); Ωrad ∼ 10−4.(Radiation ≡ relativistic particles.)
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Thermal history of the Universe
Currently: Universe dominated by dark energy (∼ 75%)and non–relativistic matter (∼ 25%); Ωrad ∼ 10−4.(Radiation ≡ relativistic particles.)
Dependence on scale factor R: ρm ∝ R−3, ρrad ∝ R−4
ρ: energy density, units GeV4
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Thermal history of the Universe
Currently: Universe dominated by dark energy (∼ 75%)and non–relativistic matter (∼ 25%); Ωrad ∼ 10−4.(Radiation ≡ relativistic particles.)
Dependence on scale factor R: ρm ∝ R−3, ρrad ∝ R−4
ρ: energy density, units GeV4
Implies ρrad > ρm for R < 5 · 10−4R0, i.e. T >∼ 1 eV
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Thermal history of the Universe
Currently: Universe dominated by dark energy (∼ 75%)and non–relativistic matter (∼ 25%); Ωrad ∼ 10−4.(Radiation ≡ relativistic particles.)
Dependence on scale factor R: ρm ∝ R−3, ρrad ∝ R−4
ρ: energy density, units GeV4
Implies ρrad > ρm for R < 5 · 10−4R0, i.e. T >∼ 1 eV
Early Universe was dominated by radiation! (Except insome extreme ‘quintessence’ or ‘brane cosmology’models.)
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Thermal DM production
Let χ be a generic DM particle, nχ its number density (unit:GeV3). Assume χ = χ, i.e. χχ ↔SM particles is possible,but single production of χ is forbidden by some symmetry.
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Thermal DM production
Let χ be a generic DM particle, nχ its number density (unit:GeV3). Assume χ = χ, i.e. χχ ↔SM particles is possible,but single production of χ is forbidden by some symmetry.
Evolution of nχ determined by Boltzmann equation:
dnχ
dt+ 3Hnχ = −〈σannv〉
(
n2χ − n2
χ, eq
)
H = R/R : Hubble parameter〈. . . 〉 : Thermal averagingσann = σ(χχ → SM particles)v : relative velocity between χ’s in their cmsnχ, eq : χ density in full equilibrium
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dnχ
dt + 3Hnχ = −〈σannv〉(
n2χ − n2
χ, eq
)
2nd lhs term: Describes χ dilution by expansion of Universe:dR−3
dt = −3R−4R = −3HR−3
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dnχ
dt + 3Hnχ = −〈σannv〉(
n2χ − n2
χ, eq
)
2nd lhs term: Describes χ dilution by expansion of Universe:dR−3
dt = −3R−4R = −3HR−3
1st rhs term: describes χ pair annihilation; assumes shapeof nχ same as that of nχ, eq: reactions χ + f ↔ χ + f arevery fast (f : some SM particle).
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dnχ
dt + 3Hnχ = −〈σannv〉(
n2χ − n2
χ, eq
)
2nd lhs term: Describes χ dilution by expansion of Universe:dR−3
dt = −3R−4R = −3HR−3
1st rhs term: describes χ pair annihilation; assumes shapeof nχ same as that of nχ, eq: reactions χ + f ↔ χ + f arevery fast (f : some SM particle).
2nd rhs term: describes χ pair production; assumes CPconservation (=⇒ same matrix element).
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dnχ
dt + 3Hnχ = −〈σannv〉(
n2χ − n2
χ, eq
)
2nd lhs term: Describes χ dilution by expansion of Universe:dR−3
dt = −3R−4R = −3HR−3
1st rhs term: describes χ pair annihilation; assumes shapeof nχ same as that of nχ, eq: reactions χ + f ↔ χ + f arevery fast (f : some SM particle).
2nd rhs term: describes χ pair production; assumes CPconservation (=⇒ same matrix element).
Check: creation and annihilation balance iff nχ = nχ, eq.
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Rewriting the Boltzmann equation
In order to get rid of the 3Hnχ term: introduceYχ ≡ nχ
s (s: entropy density)
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Rewriting the Boltzmann equation
In order to get rid of the 3Hnχ term: introduceYχ ≡ nχ
s (s: entropy density)
For adiabatic expansion of the Universe: dsdt = −3Hs
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Rewriting the Boltzmann equation
In order to get rid of the 3Hnχ term: introduceYχ ≡ nχ
s (s: entropy density)
For adiabatic expansion of the Universe: dsdt = −3Hs
dYχ
dt = 1s
dnχ
dt − nχ
s2dsdt
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Rewriting the Boltzmann equation
In order to get rid of the 3Hnχ term: introduceYχ ≡ nχ
s (s: entropy density)
For adiabatic expansion of the Universe: dsdt = −3Hs
dYχ
dt = 1s
dnχ
dt − nχ
s2dsdt
= 1s
[
−3Hnχ − 〈σannv〉(
n2χ − n2
χ, eq
)]
+ nχ
s2 3Hs
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Rewriting the Boltzmann equation
In order to get rid of the 3Hnχ term: introduceYχ ≡ nχ
s (s: entropy density)
For adiabatic expansion of the Universe: dsdt = −3Hs
dYχ
dt = 1s
dnχ
dt − nχ
s2dsdt
= 1s
[
−3Hnχ − 〈σannv〉(
n2χ − n2
χ, eq
)]
+ nχ
s2 3Hs
= −s〈σannv〉(
Y 2χ − Y 2
χ, eq.
)
s = 2π2
45 g∗T 3 (g∗ : no. of relativistic d.o.f.)
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Rewriting the Boltzmann equation
In order to get rid of the 3Hnχ term: introduceYχ ≡ nχ
s (s: entropy density)
For adiabatic expansion of the Universe: dsdt = −3Hs
dYχ
dt = 1s
dnχ
dt − nχ
s2dsdt
= 1s
[
−3Hnχ − 〈σannv〉(
n2χ − n2
χ, eq
)]
+ nχ
s2 3Hs
= −s〈σannv〉(
Y 2χ − Y 2
χ, eq.
)
s = 2π2
45 g∗T 3 (g∗ : no. of relativistic d.o.f.)
If interactions are negligible: Yχ → const., i.e. χ density inco–moving volume is unchanged
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Rewriting the Boltzmann equation (cont’d)
Write lhs entirely in terms of dimensionless quantities:introduce x = mχ
T .
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Rewriting the Boltzmann equation (cont’d)
Write lhs entirely in terms of dimensionless quantities:introduce x = mχ
T .Had: s = −3Hs
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Rewriting the Boltzmann equation (cont’d)
Write lhs entirely in terms of dimensionless quantities:introduce x = mχ
T .Had: s = −3Hs
=⇒ ddt
(
g∗T 3)
= −3H(
g∗T 3)
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Rewriting the Boltzmann equation (cont’d)
Write lhs entirely in terms of dimensionless quantities:introduce x = mχ
T .Had: s = −3Hs
=⇒ ddt
(
g∗T 3)
= −3H(
g∗T 3)
=⇒ g∗T 3 + 3g∗T 2T = −3Hg∗T 3
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Rewriting the Boltzmann equation (cont’d)
Write lhs entirely in terms of dimensionless quantities:introduce x = mχ
T .Had: s = −3Hs
=⇒ ddt
(
g∗T 3)
= −3H(
g∗T 3)
=⇒ g∗T 3 + 3g∗T 2T = −3Hg∗T 3
=⇒ T = −(
H + g∗3g∗
)
T
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Rewriting the Boltzmann equation (cont’d)
Write lhs entirely in terms of dimensionless quantities:introduce x = mχ
T .Had: s = −3Hs
=⇒ ddt
(
g∗T 3)
= −3H(
g∗T 3)
=⇒ g∗T 3 + 3g∗T 2T = −3Hg∗T 3
=⇒ T = −(
H + g∗3g∗
)
T
From now on: set g∗ = 0, since g∗ changes only slowly withtime. (Except during QCD phase transition.)
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Rewriting the Boltzmann equation (cont’d)
Write lhs entirely in terms of dimensionless quantities:introduce x = mχ
T .Had: s = −3Hs
=⇒ ddt
(
g∗T 3)
= −3H(
g∗T 3)
=⇒ g∗T 3 + 3g∗T 2T = −3Hg∗T 3
=⇒ T = −(
H + g∗3g∗
)
T
From now on: set g∗ = 0, since g∗ changes only slowly withtime. (Except during QCD phase transition.)
=⇒ x = −mχ
T 2 T = − xT T = xH
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Boltzmann equation (cont’d)
dYχ
dx = 1x
dYχ
dt
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Boltzmann equation (cont’d)
dYχ
dx = 1x
dYχ
dt
= − 1Hxs〈σannv〉
(
Y 2χ − Y 2
χ, eq
)
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Boltzmann equation (cont’d)
dYχ
dx = 1x
dYχ
dt
= − 1Hxs〈σannv〉
(
Y 2χ − Y 2
χ, eq
)
Use s = 2π2
45 g∗T 3
H =√
ρrad
3M2P
= π√90
√g∗
MPT 2 for flat, rad.–dom. Universe
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Boltzmann equation (cont’d)
dYχ
dx = 1x
dYχ
dt
= − 1Hxs〈σannv〉
(
Y 2χ − Y 2
χ, eq
)
Use s = 2π2
45 g∗T 3
H =√
ρrad
3M2P
= π√90
√g∗
MPT 2 for flat, rad.–dom. Universe
=⇒ dYχ
dx = −4π√
g∗√90
mχMP
x2 〈σannv〉(
Y 2χ − Y 2
χ, eq
)
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Boltzmann equation (cont’d)
dYχ
dx = 1x
dYχ
dt
= − 1Hxs〈σannv〉
(
Y 2χ − Y 2
χ, eq
)
Use s = 2π2
45 g∗T 3
H =√
ρrad
3M2P
= π√90
√g∗
MPT 2 for flat, rad.–dom. Universe
=⇒ dYχ
dx = −4π√
g∗√90
mχMP
x2 〈σannv〉(
Y 2χ − Y 2
χ, eq
)
For T >∼ 200 MeV: 10 <∼4π
√g∗√
90<∼ 20 (SM, MSSM)
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Condition for thermal equilibrium
nχ〈σannv〉 > H for some T !
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Condition for thermal equilibrium
nχ〈σannv〉 > H for some T !
For renormalizable interactions: easiest to satisfy forT ∼ mχ =⇒
〈σannv〉(T ≃ mχ) >1
mχMP
(See: freeze–in).
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Condition for thermal equilibrium
nχ〈σannv〉 > H for some T !
For renormalizable interactions: easiest to satisfy forT ∼ mχ =⇒
〈σannv〉(T ≃ mχ) >1
mχMP
(See: freeze–in).
For non–renormalizable interactions: easiest to satisfyat maximal temperature, T ≃ TR. (See: G)
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Condition for thermal equilibrium
nχ〈σannv〉 > H for some T !
For renormalizable interactions: easiest to satisfy forT ∼ mχ =⇒
〈σannv〉(T ≃ mχ) >1
mχMP
(See: freeze–in).
For non–renormalizable interactions: easiest to satisfyat maximal temperature, T ≃ TR. (See: G)
For TR < mχ: Easiest to satisfy for T ≃ TR (see: WIMPat low TR).
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Example 1: WIMP
Decouple (freeze out) at temperature T ≪ mχ (see below).(N.B. Means χ makes cold DM!)
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Example 1: WIMP
Decouple (freeze out) at temperature T ≪ mχ (see below).(N.B. Means χ makes cold DM!)
χ’s are non–relativistic: 2 consequences
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Example 1: WIMP
Decouple (freeze out) at temperature T ≪ mχ (see below).(N.B. Means χ makes cold DM!)
χ’s are non–relativistic: 2 consequences
1) nχ ≃ gχ
(
mχT2π
)3/2e−x
〈σannv〉 ≃ x3/2
2√
π
∫ ∞0 dv v2(σannv)e−xv2/4
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Example 1: WIMP
Decouple (freeze out) at temperature T ≪ mχ (see below).(N.B. Means χ makes cold DM!)
χ’s are non–relativistic: 2 consequences
1) nχ ≃ gχ
(
mχT2π
)3/2e−x
〈σannv〉 ≃ x3/2
2√
π
∫ ∞0 dv v2(σannv)e−xv2/4
2) Most of the time: can expand cross section in χ velocity:
σannv = a + bv2 + . . . =⇒ 〈σannv〉 = a + 6 bx + . . .
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Example 1: WIMP
Decouple (freeze out) at temperature T ≪ mχ (see below).(N.B. Means χ makes cold DM!)
χ’s are non–relativistic: 2 consequences
1) nχ ≃ gχ
(
mχT2π
)3/2e−x
〈σannv〉 ≃ x3/2
2√
π
∫ ∞0 dv v2(σannv)e−xv2/4
2) Most of the time: can expand cross section in χ velocity:
σannv = a + bv2 + . . . =⇒ 〈σannv〉 = a + 6 bx + . . .
Typically, a, b <∼α2
m2χ, α2 ∼ 10−3, unless a is suppressed by
some symmetry; e.g. for χχ → ff : a ∝ m2f .
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Case 1: Low reheat temperature
Let TR be the highest temperature of theradiation–dominated universe (after inflation).
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Case 1: Low reheat temperature
Let TR be the highest temperature of theradiation–dominated universe (after inflation).
Boundary condition: nχ(TR) = 0 (??)
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Case 1: Low reheat temperature
Let TR be the highest temperature of theradiation–dominated universe (after inflation).
Boundary condition: nχ(TR) = 0 (??)
Assume TR < mχ/20 =⇒ χ annihilation is negligible (seebelow): ignore 1st rhs term in Boltzmann equation!
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Case 1: Low reheat temperature
Let TR be the highest temperature of theradiation–dominated universe (after inflation).
Boundary condition: nχ(TR) = 0 (??)
Assume TR < mχ/20 =⇒ χ annihilation is negligible (seebelow): ignore 1st rhs term in Boltzmann equation!
=⇒ dYχ
dx =4π
√g∗√
90
mχMP
x2
(
a + 6bx
)
Y 2χ, eq.
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Case 1: Low reheat temperature
Let TR be the highest temperature of theradiation–dominated universe (after inflation).
Boundary condition: nχ(TR) = 0 (??)
Assume TR < mχ/20 =⇒ χ annihilation is negligible (seebelow): ignore 1st rhs term in Boltzmann equation!
=⇒ dYχ
dx =4π
√g∗√
90
mχMP
x2
(
a + 6bx
)
Y 2χ, eq.
=4π
√g∗√
90
mχMP
x2
(
a + 6bx
) m6χ
(2π)3x3 e−2x 452x6
(2π2)2g2∗m6
χ
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Case 1: Low reheat temperature
Let TR be the highest temperature of theradiation–dominated universe (after inflation).
Boundary condition: nχ(TR) = 0 (??)
Assume TR < mχ/20 =⇒ χ annihilation is negligible (seebelow): ignore 1st rhs term in Boltzmann equation!
=⇒ dYχ
dx =4π
√g∗√
90
mχMP
x2
(
a + 6bx
)
Y 2χ, eq.
=4π
√g∗√
90
mχMP
x2
(
a + 6bx
) m6χ
(2π)3x3 e−2x 452x6
(2π2)2g2∗m6
χ
= 452
8√
90g3/2∗ π6
g2χmχMP (ax + 6b)e−2x.
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Case 1: Low reheat temperature
Let TR be the highest temperature of theradiation–dominated universe (after inflation).
Boundary condition: nχ(TR) = 0 (??)
Assume TR < mχ/20 =⇒ χ annihilation is negligible (seebelow): ignore 1st rhs term in Boltzmann equation!
=⇒ dYχ
dx =4π
√g∗√
90
mχMP
x2
(
a + 6bx
)
Y 2χ, eq.
=4π
√g∗√
90
mχMP
x2
(
a + 6bx
) m6χ
(2π)3x3 e−2x 452x6
(2π2)2g2∗m6
χ
= 452
8√
90g3/2∗ π6
g2χmχMP (ax + 6b)e−2x.
=⇒ Yχ(x ≫ xR) =452g2
χ
8√
90g3/2∗ π6
mχMP · e−2xR[
a2
(
xR − 12
)
+ 3b]
.
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To get current Ωχh2
Saw: Yχ → Yχ,0 = const. for x ≫ xR.
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To get current Ωχh2
Saw: Yχ → Yχ,0 = const. for x ≫ xR.
=⇒ Ωχh2 = ρχ
ρcrit.h2 = nχmχ
3H20M2
P
H20
(100 kmMpc−1 sec−1)2
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To get current Ωχh2
Saw: Yχ → Yχ,0 = const. for x ≫ xR.
=⇒ Ωχh2 = ρχ
ρcrit.h2 = nχmχ
3H20M2
P
H20
(100 kmMpc−1 sec−1)2
= Yχ,0s0mχ
3M2P (100 kmMpc−1 sec−1)2
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To get current Ωχh2
Saw: Yχ → Yχ,0 = const. for x ≫ xR.
=⇒ Ωχh2 = ρχ
ρcrit.h2 = nχmχ
3H20M2
P
H20
(100 kmMpc−1 sec−1)2
= Yχ,0s0mχ
3M2P (100 kmMpc−1 sec−1)2
Use 1 Mpc = 3.09 · 1019 km, 1 sec−1 = 6.6 · 10−25 GeV,s0 = 2.9 · 103 cm−3 = 2.2 · 10−38 GeV3, and introducedimensionless quantities a = am2
χ, b = bm2
χ
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To get current Ωχh2
Saw: Yχ → Yχ,0 = const. for x ≫ xR.
=⇒ Ωχh2 = ρχ
ρcrit.h2 = nχmχ
3H20M2
P
H20
(100 kmMpc−1 sec−1)2
= Yχ,0s0mχ
3M2P (100 kmMpc−1 sec−1)2
Use 1 Mpc = 3.09 · 1019 km, 1 sec−1 = 6.6 · 10−25 GeV,s0 = 2.9 · 103 cm−3 = 2.2 · 10−38 GeV3, and introducedimensionless quantities a = am2
χ, b = bm2
χ
=⇒ Ωχh2 = mχYχ,0 2.8 · 108 GeV−1
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To get current Ωχh2
Saw: Yχ → Yχ,0 = const. for x ≫ xR.
=⇒ Ωχh2 = ρχ
ρcrit.h2 = nχmχ
3H20M2
P
H20
(100 kmMpc−1 sec−1)2
= Yχ,0s0mχ
3M2P (100 kmMpc−1 sec−1)2
Use 1 Mpc = 3.09 · 1019 km, 1 sec−1 = 6.6 · 10−25 GeV,s0 = 2.9 · 103 cm−3 = 2.2 · 10−38 GeV3, and introducedimensionless quantities a = am2
χ, b = bm2
χ
=⇒ Ωχh2 = mχYχ,0 2.8 · 108 GeV−1
=⇒ Ωχh2 = 0.9 · 1023 e−2mχ/TR
[
a2
(
mχ
TR− 1
2
)
+ 3b]
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To get current Ωχh2
Saw: Yχ → Yχ,0 = const. for x ≫ xR.
=⇒ Ωχh2 = ρχ
ρcrit.h2 = nχmχ
3H20M2
P
H20
(100 kmMpc−1 sec−1)2
= Yχ,0s0mχ
3M2P (100 kmMpc−1 sec−1)2
Use 1 Mpc = 3.09 · 1019 km, 1 sec−1 = 6.6 · 10−25 GeV,s0 = 2.9 · 103 cm−3 = 2.2 · 10−38 GeV3, and introducedimensionless quantities a = am2
χ, b = bm2
χ
=⇒ Ωχh2 = mχYχ,0 2.8 · 108 GeV−1
=⇒ Ωχh2 = 0.9 · 1023 e−2mχ/TR
[
a2
(
mχ
TR− 1
2
)
+ 3b]
Example: a = 0, b = 10−4 =⇒ need TR ≃ 0.04mχ
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Case 2: Thermal WIMP
Assume χ was in full thermal equilibrium after inflation.
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Case 2: Thermal WIMP
Assume χ was in full thermal equilibrium after inflation.
Requiresnχ〈σannv〉 > H
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Case 2: Thermal WIMP
Assume χ was in full thermal equilibrium after inflation.
Requiresnχ〈σannv〉 > H
For T < mχ : nχ ≃ nχ, eq ∝ T 3/2e−mχ/T , H ∝ T 2
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Case 2: Thermal WIMP
Assume χ was in full thermal equilibrium after inflation.
Requiresnχ〈σannv〉 > H
For T < mχ : nχ ≃ nχ, eq ∝ T 3/2e−mχ/T , H ∝ T 2
Inequality cannot be true for arbitrarily small T ; point whereinequality becomes (approximate) equality definesdecoupling (freeze–out) temperature TF .
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Case 2: Thermal WIMP
Assume χ was in full thermal equilibrium after inflation.
Requiresnχ〈σannv〉 > H
For T < mχ : nχ ≃ nχ, eq ∝ T 3/2e−mχ/T , H ∝ T 2
Inequality cannot be true for arbitrarily small T ; point whereinequality becomes (approximate) equality definesdecoupling (freeze–out) temperature TF .
For T < TF : WIMP production negligible, only annihilationrelevant in Boltzmann equation.
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Thermal WIMP: solution of Boltzmann eq.
Had dYχ
dx = −4π√
g∗√90
mχMP
x2 〈σannv〉(
Y 2χ − Y 2
χ, eq
)
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Thermal WIMP: solution of Boltzmann eq.
Had dYχ
dx = −4π√
g∗√90
mχMP
x2 〈σannv〉(
Y 2χ − Y 2
χ, eq
)
High temperature, T > TF : write Yχ = Yχ, eq. + ∆, ignore ∆2
term
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Thermal WIMP: solution of Boltzmann eq.
Had dYχ
dx = −4π√
g∗√90
mχMP
x2 〈σannv〉(
Y 2χ − Y 2
χ, eq
)
High temperature, T > TF : write Yχ = Yχ, eq. + ∆, ignore ∆2
term=⇒ d∆
dx = −dYχ, eq
dx + dYχ
dx
≃ −dYχ, eq
dx − 4π√
g∗√90
mχMP
x2 〈σannv〉 (2Yχ, eq∆)
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Thermal WIMP: solution of Boltzmann eq.
Had dYχ
dx = −4π√
g∗√90
mχMP
x2 〈σannv〉(
Y 2χ − Y 2
χ, eq
)
High temperature, T > TF : write Yχ = Yχ, eq. + ∆, ignore ∆2
term=⇒ d∆
dx = −dYχ, eq
dx + dYχ
dx
≃ −dYχ, eq
dx − 4π√
g∗√90
mχMP
x2 〈σannv〉 (2Yχ, eq∆)
Use dYχ, eq
dx = −32
Yχ, eq
x − Yχ, eq ≃ −Yχ, eq (x ≫ 1):
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Thermal WIMP: solution of Boltzmann eq.
Had dYχ
dx = −4π√
g∗√90
mχMP
x2 〈σannv〉(
Y 2χ − Y 2
χ, eq
)
High temperature, T > TF : write Yχ = Yχ, eq. + ∆, ignore ∆2
term=⇒ d∆
dx = −dYχ, eq
dx + dYχ
dx
≃ −dYχ, eq
dx − 4π√
g∗√90
mχMP
x2 〈σannv〉 (2Yχ, eq∆)
Use dYχ, eq
dx = −32
Yχ, eq
x − Yχ, eq ≃ −Yχ, eq (x ≫ 1):To keep d∆
dx = 0 : need
∆ ≃ x2
2.64mχMP√
g∗〈σannv〉
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Low−T solution
Yχ, eq → 0 =⇒ can ignore production term in Boltzmann eq.
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Low−T solution
Yχ, eq → 0 =⇒ can ignore production term in Boltzmann eq.
d∆dx ≃ −1.32mχMP
√g∗x−2〈σannv〉∆2
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Low−T solution
Yχ, eq → 0 =⇒ can ignore production term in Boltzmann eq.
d∆dx ≃ −1.32mχMP
√g∗x−2〈σannv〉∆2
=⇒ 1∆(xF ) −
1∆(∞) = −1.32mχMP
√g∗
∫ ∞xF
dxx−2〈σannv〉≡ −1.32mχMP
√g∗ J(xF )
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Low−T solution
Yχ, eq → 0 =⇒ can ignore production term in Boltzmann eq.
d∆dx ≃ −1.32mχMP
√g∗x−2〈σannv〉∆2
=⇒ 1∆(xF ) −
1∆(∞) = −1.32mχMP
√g∗
∫ ∞xF
dxx−2〈σannv〉≡ −1.32mχMP
√g∗ J(xF )
Assume ∆(∞) ≪ ∆(xF )
=⇒ Yχ,0 = 11.32
√g∗mχMP J(xF )
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Low−T solution
Yχ, eq → 0 =⇒ can ignore production term in Boltzmann eq.
d∆dx ≃ −1.32mχMP
√g∗x−2〈σannv〉∆2
=⇒ 1∆(xF ) −
1∆(∞) = −1.32mχMP
√g∗
∫ ∞xF
dxx−2〈σannv〉≡ −1.32mχMP
√g∗ J(xF )
Assume ∆(∞) ≪ ∆(xF )
=⇒ Yχ,0 = 11.32
√g∗mχMP J(xF )
=⇒ Ωχh2 ≃ 8.7·10−11 GeV−2
√g∗J(xF )
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Low−T solution
Yχ, eq → 0 =⇒ can ignore production term in Boltzmann eq.
d∆dx ≃ −1.32mχMP
√g∗x−2〈σannv〉∆2
=⇒ 1∆(xF ) −
1∆(∞) = −1.32mχMP
√g∗
∫ ∞xF
dxx−2〈σannv〉≡ −1.32mχMP
√g∗ J(xF )
Assume ∆(∞) ≪ ∆(xF )
=⇒ Yχ,0 = 11.32
√g∗mχMP J(xF )
=⇒ Ωχh2 ≃ 8.7·10−11 GeV−2
√g∗J(xF )
Typically, xF ≃ 22; depends only logarithmically on σann.
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Low−T solution
Yχ, eq → 0 =⇒ can ignore production term in Boltzmann eq.
d∆dx ≃ −1.32mχMP
√g∗x−2〈σannv〉∆2
=⇒ 1∆(xF ) −
1∆(∞) = −1.32mχMP
√g∗
∫ ∞xF
dxx−2〈σannv〉≡ −1.32mχMP
√g∗ J(xF )
Assume ∆(∞) ≪ ∆(xF )
=⇒ Yχ,0 = 11.32
√g∗mχMP J(xF )
=⇒ Ωχh2 ≃ 8.7·10−11 GeV−2
√g∗J(xF )
Typically, xF ≃ 22; depends only logarithmically on σann.
Non-relativistic expansion: J(xF ) = axF
+ 3bx2
F. . .
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Ωχh2 ≃ 8.7·10−11 GeV−2
√g∗J(xF )
Solution validated numerically.
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Ωχh2 ≃ 8.7·10−11 GeV−2
√g∗J(xF )
Solution validated numerically.
Density has no explicit dependence on mχ.
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Ωχh2 ≃ 8.7·10−11 GeV−2
√g∗J(xF )
Solution validated numerically.
Density has no explicit dependence on mχ.
Density has no dependence on reheat temperature TR,if TR > TF .
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Ωχh2 ≃ 8.7·10−11 GeV−2
√g∗J(xF )
Solution validated numerically.
Density has no explicit dependence on mχ.
Density has no dependence on reheat temperature TR,if TR > TF .
Density scales like inverse of annihilation cross section:The stronger the WIMPs annihilate, the fewer are left.
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Ωχh2 ≃ 8.7·10−11 GeV−2
√g∗J(xF )
Solution validated numerically.
Density has no explicit dependence on mχ.
Density has no dependence on reheat temperature TR,if TR > TF .
Density scales like inverse of annihilation cross section:The stronger the WIMPs annihilate, the fewer are left.
Smooth transition to previous case (TR < TF ): MD,Iminniyaz, Kakizaki, hep-ph/0603165
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Co–annihilation
Is important for SUSY scenarios with small mass splittingbetween LSP and NLSP: δm ≡ mχ′ − mχ ≪ mχ
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Co–annihilation
Is important for SUSY scenarios with small mass splittingbetween LSP and NLSP: δm ≡ mχ′ − mχ ≪ mχ
Rate (χ + f ↔ χ′ + f ′) ≫ Rate(χχ ↔ ff ), by factor∝ e(2mχ−mχ′)/T (f, f ′ : SM particles)
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Co–annihilation
Is important for SUSY scenarios with small mass splittingbetween LSP and NLSP: δm ≡ mχ′ − mχ ≪ mχ
Rate (χ + f ↔ χ′ + f ′) ≫ Rate(χχ ↔ ff ), by factor∝ e(2mχ−mχ′)/T (f, f ′ : SM particles)
χ, χ′ retain relative equilibrium well after sparticlesdecouple from SM particles: nχ′ = nχe−δm/T
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Co–annihilation
Is important for SUSY scenarios with small mass splittingbetween LSP and NLSP: δm ≡ mχ′ − mχ ≪ mχ
Rate (χ + f ↔ χ′ + f ′) ≫ Rate(χχ ↔ ff ), by factor∝ e(2mχ−mχ′)/T (f, f ′ : SM particles)
χ, χ′ retain relative equilibrium well after sparticlesdecouple from SM particles: nχ′ = nχe−δm/T
Previous treatment still applies, with replacement:
σann → σeff ∼ σann + fBσ(χχ′ → SM) + f2Bσ(χ′χ′ → SM)
fB : relative Boltzmann factor =(
1 + δmmχ
)3/2e−δm/T
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Co–annihilation
Is important for SUSY scenarios with small mass splittingbetween LSP and NLSP: δm ≡ mχ′ − mχ ≪ mχ
Rate (χ + f ↔ χ′ + f ′) ≫ Rate(χχ ↔ ff ), by factor∝ e(2mχ−mχ′)/T (f, f ′ : SM particles)
χ, χ′ retain relative equilibrium well after sparticlesdecouple from SM particles: nχ′ = nχe−δm/T
Previous treatment still applies, with replacement:
σann → σeff ∼ σann + fBσ(χχ′ → SM) + f2Bσ(χ′χ′ → SM)
fB : relative Boltzmann factor =(
1 + δmmχ
)3/2e−δm/T
σ(χχ′), σ(χ′χ′) ≫ σ(χχ) possible!
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Case 3: Freeze–in
Hall et al., arXiv:0911.1120
Assume very weak, renormalizable interaction (FeeblyInteracting Massive Particle, FIMP): never achievedthermal equilibrium
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Case 3: Freeze–in
Hall et al., arXiv:0911.1120
Assume very weak, renormalizable interaction (FeeblyInteracting Massive Particle, FIMP): never achievedthermal equilibrium
χ pair production dominated by reactions at T ∼ mχ:independent of TR as long as TR ≫ mχ
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Case 3: Freeze–in
Hall et al., arXiv:0911.1120
Assume very weak, renormalizable interaction (FeeblyInteracting Massive Particle, FIMP): never achievedthermal equilibrium
χ pair production dominated by reactions at T ∼ mχ:independent of TR as long as TR ≫ mχ
Final relic density proportional to cross section,independent of FIMP mass
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Thermal WIMPs, FIMPs: Assumptions
χ is effectively stable, τχ ≫ τU: partly testable atcolliders
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Thermal WIMPs, FIMPs: Assumptions
χ is effectively stable, τχ ≫ τU: partly testable atcolliders
No entropy production after χ decoupled: Not testableat colliders
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Thermal WIMPs, FIMPs: Assumptions
χ is effectively stable, τχ ≫ τU: partly testable atcolliders
No entropy production after χ decoupled: Not testableat colliders
H at time of χ decoupling is known: partly testable atcolliders
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Thermal Gravitino Dark Matter
Each gravitino coupling gives factor msparticlesmGMP
in cross
section, if mG ≪ √s, msparticle
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Thermal Gravitino Dark Matter
Each gravitino coupling gives factor msparticlesmGMP
in cross
section, if mG ≪ √s, msparticle
=⇒ Most important G production mechanism for mG>∼
MeV: associated production with other sparticle!
σG ≃ 124π(mGMP )2
(
26g2sM
2g + . . .
)
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Thermal Gravitino Dark Matter
Each gravitino coupling gives factor msparticlesmGMP
in cross
section, if mG ≪ √s, msparticle
=⇒ Most important G production mechanism for mG>∼
MeV: associated production with other sparticle!
σG ≃ 124π(mGMP )2
(
26g2sM
2g + . . .
)
G annihilation can be ignored; write Boltzmann eq. forYG ≡ nG/nγ:
dYG
dT = − nγσG
4TH(T )
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Gravitino DM (cont.’d)
Solution of Boltzmann eq.:
YG,0 =4nγ(TR)σG
4H(TR) ∝ TR (assuming YG(TR) = 0)
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Gravitino DM (cont.’d)
Solution of Boltzmann eq.:
YG,0 =4nγ(TR)σG
4H(TR) ∝ TR (assuming YG(TR) = 0)
=⇒ ΩGh2 ≃ 0.1(
Mg
1 TeV
)21 GeV
mG
TR
2.4·107 GeV
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Gravitino DM (cont.’d)
Solution of Boltzmann eq.:
YG,0 =4nγ(TR)σG
4H(TR) ∝ TR (assuming YG(TR) = 0)
=⇒ ΩGh2 ≃ 0.1(
Mg
1 TeV
)21 GeV
mG
TR
2.4·107 GeV
Inclusion of thermal corrections: e.g. Pradler & Steffen,hep-ph/0612291
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Gravitino DM (cont.’d)
Solution of Boltzmann eq.:
YG,0 =4nγ(TR)σG
4H(TR) ∝ TR (assuming YG(TR) = 0)
=⇒ ΩGh2 ≃ 0.1(
Mg
1 TeV
)21 GeV
mG
TR
2.4·107 GeV
Inclusion of thermal corrections: e.g. Pradler & Steffen,hep-ph/0612291In general, have to add ΩNLSP
mG
mNLSPfrom (late) decays of
NLSPs. (BBN!)
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DM Production from Inflaton Decay
Only consider perturbative decays here.
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DM Production from Inflaton Decay
Only consider perturbative decays here.
χ : DM particle; φ : inflaton
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DM Production from Inflaton Decay
Only consider perturbative decays here.
χ : DM particle; φ : inflaton
B(φ → χ) : Average number of χ particles produced per φdecay
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DM Production from Inflaton Decay
Only consider perturbative decays here.
χ : DM particle; φ : inflaton
B(φ → χ) : Average number of χ particles produced per φdecay
Instantaneous φ decay approximation: all inflatons decay atT = TR.
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DM Production from Inflaton Decay
Only consider perturbative decays here.
χ : DM particle; φ : inflaton
B(φ → χ) : Average number of χ particles produced per φdecay
Instantaneous φ decay approximation: all inflatons decay atT = TR.
Inflatons are non–relativistic when they decay.
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DM Production from Inflaton Decay (cont.’d)
Energy conserved during φ decay
=⇒ nφmφ = ρrad(TR) = π2
30g∗T 4R
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DM Production from Inflaton Decay (cont.’d)
Energy conserved during φ decay
=⇒ nφmφ = ρrad(TR) = π2
30g∗T 4R
=⇒ Yχ(TR) = nχ(TR)s(TR) = B(φ→χ)nφ(TR)
s(TR)
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DM Production from Inflaton Decay (cont.’d)
Energy conserved during φ decay
=⇒ nφmφ = ρrad(TR) = π2
30g∗T 4R
=⇒ Yχ(TR) = nχ(TR)s(TR) = B(φ→χ)nφ(TR)
s(TR)
= B(φ→χ)ρrad(TR)mφs(TR)
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DM Production from Inflaton Decay (cont.’d)
Energy conserved during φ decay
=⇒ nφmφ = ρrad(TR) = π2
30g∗T 4R
=⇒ Yχ(TR) = nχ(TR)s(TR) = B(φ→χ)nφ(TR)
s(TR)
= B(φ→χ)ρrad(TR)mφs(TR)
= 34
TR
mφB(φ → χ)
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DM Production from Inflaton Decay (cont.’d)
Energy conserved during φ decay
=⇒ nφmφ = ρrad(TR) = π2
30g∗T 4R
=⇒ Yχ(TR) = nχ(TR)s(TR) = B(φ→χ)nφ(TR)
s(TR)
= B(φ→χ)ρrad(TR)mφs(TR)
= 34
TR
mφB(φ → χ)
If χ production and annihilation at T < TR is negligible,universe evolves adiabatically:
=⇒ Ωχh2 = 2.1 · 108 mχ
mφ
TR
1 GeVB(φ → χ)
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Possibilities forB(φ → χ)
If χ = LSP: expect B(φ → χ) ≃ 1: Excludes chargedLSP for mφ > 2mχ, TR >∼ 1 MeV!
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Possibilities forB(φ → χ)
If χ = LSP: expect B(φ → χ) ≃ 1: Excludes chargedLSP for mφ > 2mχ, TR >∼ 1 MeV!
“Democratic” coupling: B(φ → χ) ≃ gχ/g∗ ∼ 10−2.
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Possibilities forB(φ → χ)
If χ = LSP: expect B(φ → χ) ≃ 1: Excludes chargedLSP for mφ > 2mχ, TR >∼ 1 MeV!
“Democratic” coupling: B(φ → χ) ≃ gχ/g∗ ∼ 10−2.
φ → ffχχ (4–body):
B(φ → χ) ∼ α2χ
96π3
(
1 − 4m2χ
m2φ
)2 (
1 − 2mχ
mφ
)5/2
(Assumes σ(χχ ↔ ff) ∼ α2χ
m2χ, φ → ff dominates.)
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Possibilities forB(φ → χ)
If χ = LSP: expect B(φ → χ) ≃ 1: Excludes chargedLSP for mφ > 2mχ, TR >∼ 1 MeV!
“Democratic” coupling: B(φ → χ) ≃ gχ/g∗ ∼ 10−2.
φ → ffχχ (4–body):
B(φ → χ) ∼ α2χ
96π3
(
1 − 4m2χ
m2φ
)2 (
1 − 2mχ
mφ
)5/2
(Assumes σ(χχ ↔ ff) ∼ α2χ
m2χ, φ → ff dominates.)
Can be most important production mechanism forsuperheavy Dark Matter (mχ ∼ 1012 GeV) in chaoticinflation (mφ ∼ 1013 GeV); for LSP if TR <∼ 0.03mχ; . . .
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Summary
Different production mechanisms give very differentresults for Ωχh2:
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Summary
Different production mechanisms give very differentresults for Ωχh2:
Thermal WIMP: Ωχh2 ∝ 1〈σeffv〉 , independent of TR: most
frequently studied case; needs TR >∼ 0.05mχ.
Making Dark Matter – p. 35/35
![Page 169: Making Dark Matter - uni-bonn.de · Making Dark Matter Manuel Drees Bonn University & Bethe Center for Theoretical Physics Making Dark Matter – p. 1/35](https://reader033.vdocument.in/reader033/viewer/2022041918/5e6afaec058f43024b113d42/html5/thumbnails/169.jpg)
Summary
Different production mechanisms give very differentresults for Ωχh2:
Thermal WIMP: Ωχh2 ∝ 1〈σeffv〉 , independent of TR: most
frequently studied case; needs TR >∼ 0.05mχ.
WIMP that never was in equilibrium:Ωχh2 ∝ e−2mχ/TRm2
χ〈σannv〉
Making Dark Matter – p. 35/35
![Page 170: Making Dark Matter - uni-bonn.de · Making Dark Matter Manuel Drees Bonn University & Bethe Center for Theoretical Physics Making Dark Matter – p. 1/35](https://reader033.vdocument.in/reader033/viewer/2022041918/5e6afaec058f43024b113d42/html5/thumbnails/170.jpg)
Summary
Different production mechanisms give very differentresults for Ωχh2:
Thermal WIMP: Ωχh2 ∝ 1〈σeffv〉 , independent of TR: most
frequently studied case; needs TR >∼ 0.05mχ.
WIMP that never was in equilibrium:Ωχh2 ∝ e−2mχ/TRm2
χ〈σannv〉
Thermal gravitino production: ΩGh2 ∝ TR
mG.
Making Dark Matter – p. 35/35
![Page 171: Making Dark Matter - uni-bonn.de · Making Dark Matter Manuel Drees Bonn University & Bethe Center for Theoretical Physics Making Dark Matter – p. 1/35](https://reader033.vdocument.in/reader033/viewer/2022041918/5e6afaec058f43024b113d42/html5/thumbnails/171.jpg)
Summary
Different production mechanisms give very differentresults for Ωχh2:
Thermal WIMP: Ωχh2 ∝ 1〈σeffv〉 , independent of TR: most
frequently studied case; needs TR >∼ 0.05mχ.
WIMP that never was in equilibrium:Ωχh2 ∝ e−2mχ/TRm2
χ〈σannv〉
Thermal gravitino production: ΩGh2 ∝ TR
mG.
Production from inflaton decay: Ωχh2 ∝ mχTRB(φ→χ)mφ
.
Making Dark Matter – p. 35/35
![Page 172: Making Dark Matter - uni-bonn.de · Making Dark Matter Manuel Drees Bonn University & Bethe Center for Theoretical Physics Making Dark Matter – p. 1/35](https://reader033.vdocument.in/reader033/viewer/2022041918/5e6afaec058f43024b113d42/html5/thumbnails/172.jpg)
Summary
Different production mechanisms give very differentresults for Ωχh2:
Thermal WIMP: Ωχh2 ∝ 1〈σeffv〉 , independent of TR: most
frequently studied case; needs TR >∼ 0.05mχ.
WIMP that never was in equilibrium:Ωχh2 ∝ e−2mχ/TRm2
χ〈σannv〉
Thermal gravitino production: ΩGh2 ∝ TR
mG.
Production from inflaton decay: Ωχh2 ∝ mχTRB(φ→χ)mφ
.
Only the thermal WIMP scenario can be tested usingcollider data and results from WIMP searchexperiments. Other scenarios can only be tested withadditional input to constrain cosmology (TR, . . . ).
Making Dark Matter – p. 35/35