A brief thermal history of the UniverseA brief thermal history of the Universe Martin Kunz...
Transcript of A brief thermal history of the UniverseA brief thermal history of the Universe Martin Kunz...
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A brief thermal history of the Universe
Martin Kunz Université de Genève
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Overview
• Equilibrium description – Distribution functions – Neutrino decoupling and the neutrino
background – Photon decoupling, recombination and the CMB – BBN and light element abundances – Comparison to observations
• Basics of Boltzmann equation – hot and cold relics
• Summary
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Brief history of the Universe
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Equilibrium distributions Short-range interactions maintaining
thermodynamic equilibrium: , T temperature, µ chem. pot.
f(k, t)d3k =
g
(2�)
3
�exp[(E � µ)/T ]± 1
��1d3k
E =p
k2 + m2
n =Z
f(k)d3k
� =Z
E(k)f(k)d3k
p =Z |k|2
3E(k)f(k)d3k
number density
energy density
pressure
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Relativistic species, m << T
Crank handle, using m = µ = 0 -> E ~ k (use x=E/T as integration variable)
nB = T 3 g�(3)⇥2
nF =34nB
⇥B = T 4 g
30�2 ⇥F =
78⇥B
p =�
3
with ργ ~ a-4 => Tγ ~ 1/a -> expanding universe cools down
-> wrad = prad/ρrad = 1/3
-> Stefan-Boltzmann law ργ ~ T4
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Massive species, m >> T
Expand and neglect +/-1 wrt exp(m/T)
E =p
k2 + m2 = mp
1 + k2/m2 ⇡ m + k2/(2m)
n = g
✓mT
2�
◆3/2
e�(m�µ)/T
� = mn +32nT
p = nT ⌧ �
Ekin =32kBT-> Ekin / particle:
Massive par*cles are suppressed by Boltzmann factor exp(-‐m/T), so they will quickly drop out of thermal equilibrium when T < m -‐> ‘freeze out’ -‐> effec*ve μ
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Multiple relativistic species If we have several species at different temperatures:
⇥R =T 4
�
30�2g⇤ g⇤ =
X
i2B
gi
✓Ti
T�
◆4
+78
X
j2F
gj
✓Tj
T�
◆4
Entropy density: s =� + p
T/ T 3 d(sa3)/dt = 0
(use f and ) � + 3H(� + p) = 0
s =2�2
45g⇤ST 3
� g⇤S =X
i2B
gi
✓Ti
T�
◆3
+78
X
j2F
gj
✓Tj
T�
◆3
• T� / g�1/3⇤S a�1
Now we are ready to study par.cle evolu.on in the early universe!
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Neutrino decoupling
Interaction rate: Γ species in equil.: Γ >> H Expansion rate: H species decoupled: Γ << H
�(T ) = n(T )h�viT �F ' G2F E2 ' G2
F T 2 �F ⇠ G2F T 5
H(T ) =r
8�G
3p
⇥R '5.44mP
T 2 g⇤ = 2 +78(3⇥ 2 + 2⇥ 2)
) �F
H(T )' 0.24T 3G2
F mP '✓
T
1MeV
◆3
Neutrinos decouple when temperature drops below ~ 1 MeV because their interactions become too weak.
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Temperature of ν background
Shortly after the neutrinos decouple, we reach T=0.5MeV=me and the entropy in electron-positron pairs is transferred to photons but not to the neutrinos. Photon + electron entropy g*S(Ta)3 is separately conserved:
g⇤(T�dec > T > me) = 2 +78⇥ 4 =
112
, g⇤(T < me) = 2
How much are the photons heated by the electron-positron annihilation?
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Temperature of ν background
Shortly after the neutrinos decouple, we reach T=0.5MeV=me and the entropy in electron-positron pairs is transferred to photons but not to the neutrinos. Photon + electron entropy g*S(Ta)3 is separately conserved:
g⇤(T�dec > T > me) = 2 +78⇥ 4 =
112
, g⇤(T < me) = 2
(aT�)3after
(aT�)3before
=(g⇤)before
(g⇤)after=
114
Since (aTν) = (aTγ)before we now have Tγ = (11/4)1/3 Tν -> for T<0.5me : g* ~ 3.36 and g*S ~ 3.91 for radiation (γ+ν)
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Photon decoupling
e+/e- annihilation stopped by baryon asymmetry, remaining electrons and photons stay in thermal contact (Compton scattering) until electrons and protons form neutral hydrogen (recombination). As number of free electrons ne drops, photons decouple when Γγ ~ H.
�� = ne⇤T , ⇤T =8⇥�2
EM
3m2e
nj = gj
✓mjT
2�
◆3/2
e(µj�mj)/T , j = e, p,H
µp + µe = µH
mp + me = mH + B
gp = ge = 2, gH = 4
due to interactions (number conservation) -> use this to remove µ’s
energy conservation, binding energy: B=13.6eV
nB = np + nH = ne + nH (np = ne)Baryon number density:
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Photon decoupling II We can write nH
nenp⇡ gH
gpge
✓meT
2�
◆�3/2
eB/T
introduce Xe=ne/nB (fractional ionisation) and notice that (in equilibrium)
1�Xe
X2e
=nHnB
nenp=
4p
2�(3)p⇥
nB
n�
✓T
me
◆3/2
eB/T Saha eqn.
Assume recombination ~ Xe=0.1 and with η ~ 10-10 [why?] => Trec ~ 0.31 eV, zrec ~ 1300 (why Trec << B?) Go back to Γγ and compare to matter dominated expansion H=H0Ωm(1+z)3/2
⇒ Tdec ~ 0.26 eV, zdec ~ 1100 -> origin of CMB!
Notice that recombination and photon decoupling are two different processes, although they happen at nearly the same time. Photons decouple because ne drops due to recombination (else zdec ~ 40!).
η
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‘equilibrium’ BBN
• T > 1MeV: p and n in equilibrium through weak interactions • T ~ 1MeV: weak interactions too slow, ν freeze-out • Light element binding energies: a few MeV -> when is it favourable to create the light elements? same game as before: for species with mass number A = #n
+#p and charge Z = #p, assumed in equilibrium with p & n
nA = gA
✓mAT
2�
◆3/2
e(µA�mA)/T = gAA3/2
2A
✓mNT
2�
◆3(1�A)/2
nZp nA�Z
n eBA/T
where we used again µA=Zµp+(A-Z) µn. With XA= nA A/nN mass fraction:
XA = . . . = (const)
✓T
mN
◆3(A�1)/2
XZp XA�Z
n �A�1eBA/T
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BBN II: NSE
-> System of equations for ‘nuclear statistical equilibrium’:
1 = Xn + Xp + X2 + X3 + . . .
Xn/Xp = e�Q/T
X2 = C
✓T
mN
◆3/2
XpXn�eB2/T
X3 = . . .
etc … • Q = 1.293 MeV • B[2H] = 2.22 MeV • B[3H] = 6.92 MeV • B[3He] = 7.72 MeV • B[4He] = 28.3 MeV
(Kolb & Turner)
log(
X(e
q))
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BBN III: actual BBN
In reality the reactions drop out of equilibrium eventually, and one needs to use the Boltzmann equation. Results:
• T ~ 10 MeV+ : equilibrium, Xn=Xp=1/2, rest X << 1 • T ~ 1 MeV: n<->p freeze-out, Xn ~ 0.15, Xp ~ 0.85, NSE
okay for rest (with X << 1) • T ~ 0.1 MeV: neutrons decay, n/p~1/8, NSE breaks down
because 4He needs Deuterium (2H) which is delayed until 0.07 MeV because of high η and low B2.
• T ~ 65 keV: now synthesis of 4He can proceed, gets nearly all neutrons that are left:
Rest is hydrogen, with some traces of 2H, 3He, 7Li and 7Be.
X4He =4n4
nN= 4
nn/2nn + np
= 2nn/np
1 + nn/np= 2
1/81 + 1/8
⇡ 0.22
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(figures from A. Weiss, Einstein Online Vol. 2 (2006), 1018)
0.1 MeV 0.01 MeV
(Dodelson)
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Timeline summary
Energy (γ) *me event
1 MeV 7s neutrino freeze-‐out
0.5 MeV 10s e+/e-‐ annihila.on, Tγ ~ 1.4 Tν
70 keV 3 minutes BBN, light elements formed
0.77 eV 70’000 yr onset of maMer domina.on
0.31 eV 300’000 yr recombina.on
0.26 eV 380’000 yr photon decoupling, origin of CMB
0.2 meV 14 Gyr today
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Comparison to observations
3 classical ‘pillars’ of big-bang model: 1. Hubble law -> Monday 2. Cosmic Microwave Background 3. BBN and element abundances CMB: we expect an isotropic radiation with
thermal spectrum to fill the universe
(COBE / NASA)
Nobel prizes: 1978 CMB discovery 2006 CMB proper.es 2011 accelerated exp.
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Comparison II : BBN
BBN tests: • baryon/photon ratio η • effective # of relativistic degrees of freedom • consistency of different abundances
The CMB anisotropies (Valeria) also depend on the baryon abundance (even-odd peak heights). Results are consistent with BBN! Ωb ≈ 0.05 (what is all the rest??!!)
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Non-equilibrium treatment Looking at equilibrium quantities is very useful, but: • when / how does freeze-out really happen? • what to do when NSE breaks down • … -> Treatment with Boltzmann equation L[f] = C[f]/2
L: Liouville operator ~ d/dt C: collision operator (à particle physics!)
relativistic / covariant form: f=f(E,t) -> only α = 0 relevant, p0=E and Γ from FRW metric =>
L = p↵ �
�x↵� �↵
��p�p� �
�p↵
L[f ] =E
�
�t� a
ap2 �
�E
�f(E, t)
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Liouville operator
We only want to know the abundance -> integrate L[f]/E over
3-momentum p to get n: • first term is just dn/dt • second term: rewrite in terms of p and integrate by parts
(d/dE = E/p d/dp)
L[f ] =E
�
�t� a
ap2 �
�E
�f(E, t)
g
2�2
Zdpp2 L[f ]
E= n + 3Hn
1. no collisions: L[f]=0 -> n ~ a-3 as expected! 2. deviations from full equilibrium will be encoded in µ
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Collision operator Roughly: (# of particles ‘in’) – (# of particles ‘out’) of phase space volume d3p d3x
For (reversible) scattering of type 1+2 <-> 3+4 and with
g1
(2⇥)3
Zd3p1
C1[f ]2E1
=Z
d�1d�2d�3d�4(2⇥)4�(4)(p1 + p2 � p3 � p4)|M|2
[f3f4(1± f1)(1± f2)� f1f2(1± f3)(1± f4)]
d�i ⌘gi
(2�)3d3pi
2Ei
f(E) ⇡ e(µ�E)/Tkinetic equilibrium & low temperature: [. . .]! e�(E1+E2)/T
he(µ1+µ2)/T � e(µ3+µ4)/T
i
ni = eµi/T n(eq)i , n(eq)
i =gi
(2�)3
Zd3p e�Ei/T
⇤⇤v⌅ ⇥ 1
n(eq)1 n(eq)
2
Zd�1d�2d�3d�4e
�(E1+E2)/T (2⇥)4�(4)(p1 + p2 � p3 � p4)|M|2
with and
) n1 + 3Hn1 = n(eq)1 n(eq)
2 h�vi
n3n4
n(eq)3 n(eq)
4
� n1n2
n(eq)1 n(eq)
2
!
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Annihilation and freeze-out
Consider annihilation processes: assume Y in thermal equilibrium 1) <σv> large -> n -> n(eq)
2) <σv> small -> n ~ a-3
Introduce x=m/T and Y=n/T3 (Y~n/s, constant for passive evol.) some algebra… => freeze-out governed by Γ/H (Γ = n(eq)<σv>)
A + A$ Y + Y
nA + 3HnA = h�vi⇣(n(eq)
A )2 � n2A
⌘
x
Y(eq)A
Y 0A = � �A
H(x)
2
4
YA
Y(eq)A
!2
� 1
3
5
(with Y(eq) = 0.09 g (for fermions) if x<<1 and Y(eq) = 0.16 g x3/2 e-x if x>>1)
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Hot and cold relics Hot relics: freeze-out when still relativistic (xf<1) ->
YA(x!1) = Y
(eq)A (xf ) = 0.278gA/g⇤S(xf )
Cold relics: freeze out when xf >> 1 => YA suppressed by e-m/T
Abundance generically proportional to 1/σ We can compute ρA,0 = mA nA,0 ΩA = ρA,0/ρcrit numerically, weak cross- sections lead to Ω ~ 1 -> WIMP miracle
(Kolb & Turner)
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Abundance estimate
• Assume roughly constant ->
• Y(eq) << Y from freeze-out onwards ->
• Integrate from freeze-out to late times ->
• Usually Y∞ << Yf ->
• This depends only linearly on xf, so take e.g. xf = 10
• particle density scales like a3 after freeze out
⇒ (extra factor ~30 from entropy generation in later annihilation processes)
⇒ insert Y∞ … final:
dY
dx= � �
x2
⇣Y 2 � Y (eq)2
⌘
dY/dx ' ��Y 2/x2
1/Y1 � 1/Yf = �/xf
Y1 ⇡ xf/�
� = mn0 = mn1
✓a1
a0
◆3
= mY1T 30
✓a1T1
a0T0
◆3
' mY1T 30
30
�X ⇠⇣xf
10
⌘ ✓g⇤(m)100
◆1/2 10�39cm2
h�vi
� =m3h⇥viH(m)
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reasons for decoupling / freeze-out
• neutrinos decouple from thermal equilibrium because interactions become too weak (~ T5 scaling of interaction rate)
• photons (CMB) decouple from equilibrium because e- disappear (recombination)
• baryons freeze-out from annihilation because of baryon-antibaryon asymmetry (no more antibaryons to annihilate with)
• WIMP’s (dark matter) freeze out from because their density drops due to e-m/T Boltzmann factor (if they are WIMP’s)
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Summary
• Methods – distribution function f(t,x,v) – conservation of entropy – full / kinetic equilibrium – Boltzmann equation (evolution of f)
• Results – evolution of particle number, pressure and energy in
equilibrium – T ~ 1/a (except when g*S changes, e.g. particle annihil.) – thermal history, freeze-out of particles when interactions
become too slow, WIMPs – origin of CMB – abundances of light elements
observa*onal “pillars” of the cosmological standard model