FORMATION OF DARK GALAXIES 1812.07000 (JCAP 03 (2019) 036) Daniel - - PowerPoint PPT Presentation

FORMATION OF DARK GALAXIES

Daniel Egana-Ugrinovic

Perimeter Institute

Jae Hyeok-Chang Rouven Essig Stony Brook University Chris Kouvaris CP3-Origins

This talk is carbon neutral. www.cooleffect.org

1812.07000 (JCAP 03 (2019) 036)

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A pessimistic scenario:

Dark matter does not have any detectable non-gravitational interactions with the SM

Liu Bolin

If the dark sector interacts only gravitationally with us… …what can we learn from its particle nature?

A COMPLETELY DARK, DARK SECTOR

Progress can be made based uniquely in astronomical and cosmological observations. High precision astronomical observatories (LSST, GAIA,LIGO, etc.) will test the behavior of DM on small scales.

How do we turn this experimental program into a dark matter theory program?

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SM FORMS STARS DUE TO DISSIPATION

SM forms compact objects since:

• 1. It has self-interactions
• 2. The baryonic gas can dissipate energy (can cool)

Properties of these objects (how did they form, sizes, masses) gives information on particle interactions.

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Microscopic Properties

ASTRONOMICAL PROPERTIES FIXED BY LAGRANGIAN PARAMETERS

Example: Chandrasekhar limit for White dwarf

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MC ∼ M 3

Pl

m2

p

In principle, it is possible to obtain a map between astronomical properties and lagrangian parameters

CAN DM FORM GALAXIES AND COMPACT OBJECTS?

Halo and star formation is a complex problem in the SM. Chemistry, multiple cooling rates… Example: the SM cooling rate functions

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Bremsstrahlung, collisional excitation, ionization, recombination… All dependent on metallicity…

Sutherland & Dopita, ApJS, 88, 253

brems only

• coll. exc.

(atom-e coll) O, C, N H, He Ne, Fe, Si Molecules (only H2 in primordial gas )

OUTLINE

• 1. Present the simplest dark-sector model that has cooling and self-

interactions

• 2. Discuss the initial, linear evolution of dark-sector perturbations

starting from the primordial power spectrum.

• 3. Continue into the non-linear regime, and discuss galactic evolution

and the formation of exotic compact objects, or “dark stars”.

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A complete history of structure formation in a dissipative dark-sector

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Simplified Models for Dark-Sector Astronomy

The Standard Model Our talk

A MODEL WITH ONLY TWO PARTICLES

Dark-electron dark-photon model

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meD , mγD , αD i ¯ ΨeDγµDµΨeD − meD ¯ ΨeDΨeD − 1 4FµνF µν + m2

γDAµAµ

e−

D, e+ D, γD

This model has only three parameters

COSMOLOGICAL ABUNDANCES

In general, all species may have a significant cosmological abundance. Dark sector is asymmetric:

• 1. No annihilations within a compact object.
• 2. Avoids complications of bound states.

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symmetric asymmetric CDM Baryons

eD eD

γD

How to generate the asymmetry? Petraki, Pearce, Kusenko 1403.1077

SYMMETRIC PART DEPLETED BY ANNIHILATIONS

The symmetric part is depleted by annihilations

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symmetric asymmetric CDM Baryons

eD

γD

eD

αD ≥ 4.6 × 10−7  meD 1 MeV 10−2 f 1/2 TeD|eD dec TSM|eD dec 1/2

Condition for efficient depletion

• f symmetric part

e+

D

e−

D

γD

DARK SECTOR COLD FEW DARK PHOTONS

Dark photons may lead to overclosure or large

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asymmetric CDM Baryons

γD

TγD|γD dec ≤ 0.2 g∗S|γD dec 10 1/31 keV mγD 1/3 TSM|γD dec

Overclosure bound

eD

TγD|BBN ≤ 0.5 TSM|BBN

∆Neff

∆Neff

(see also DAO, Cyr-Racyne et.al.1310.3278)

ργD ∼ mγDT 3

γD

Dark photon matter density

MATTER BUDGET OF OUR MODEL

Asymmetric part is a small component of matter: no bounds form bullet cluster/halo shapes

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asymmetric CDM Baryons

Katz et. al. 1303.1521

f ≡ ρeD ρDM ≤ 10%

ONLY THREE PROCESSES MATTER

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Dark-electron self-interactions Compton scattering Bremsstrahlung

DARK ELECTRON THERMODYNAMICS

Even within this simple model there are three thermodynamic regimes

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• σC ∝

α2 m2

eD

σM ∝ α2 m2

eD

m4

γD

zeq = 3400

`mfp vs. LMW , `mfp = 1/(n)

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Objective Study the formation of dark-electron galaxies and their substructure We will concentrate on the formation and evolution of the dark electron galaxy within our Milky Way

STRUCTURE FORMATION HAS TWO MAIN STAGES

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• 1. Linear growth of matter overdensities
• 2. Non-linear evolution of

the resulting(dark) matter clumps

δ = δρ ρ < 1

LINEAR GROWTH OF PERTURBATIONS

Initial conditions: Harrison-Zeldovich power spectrum Linear evolution of perturbations

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⌦ δkδ∗

k

↵ ≡ (2π)3P(k) = Ak

∂2

t δk(t) + 2H∂t δk(t) +

⇥ c2

sk2/a2 − 4πGρ0

⇤ δk = 0

cs = s Te me + 4παne m2

em2 γ

Kouvaris et.al. 1507.00959

Matter overdensities grow only

• n scales larger than the Jeans length

JEANS CRITERION DECIDES WHICH PERTURBATIONS GROW

The Jeans criterion is


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δρ ρ ∝ a

M > mJ = π 6 c3

s

✓ π ρ0(z)G ◆3/2 ρeD,

ONLY IN PARTS OF PARAMETER SPACE A MW CAN BE FORMED

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• SM

White: regions of parameter space where a Milky Way-sized perturbation may grow after equality

σC ∝ α2 m2

eD

GALAXY GOES NON-LINEAR AT

At some point, perturbations become non-linear Gravitational pull overcomes Hubble expansion: perturbations “turn-around” The galaxy’s turnaround redshift can be estimated by 


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z ≈ 2

k3

MW

2π2 P(kMW, zta) = 1 → zta ≈ 2

δρ ρ ∼ 1

tff ∼ H−1

tff ≡ ✓ 1 16πGρ ◆1/2

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Primordial regions of under- and over-densities Density contrast grows as Nonlinearities and turnaround

tff ∼ H−1

tff = ✓ 3π 32Gmene ◆1/2

Nonlinear regime: self-gravitating gas decoupled from Hubble flow

δρ ρ ∼ a δρ ρ

• δρ

ρ ∼ 1

Summary of linear perturbation growth

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Primordial regions of under- and over-densities Density contrast grows as Nonlinearities and turnaround

tff ∼ H−1

tff = ✓ 3π 32Gmene ◆1/2

Nonlinear regime: self-gravitating gas decoupled from Hubble flow

δρ ρ ∼ a δρ ρ

• δρ

ρ ∼ 1

Summary of linear perturbation growth

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Primordial regions of under- and over-densities Density contrast grows as Nonlinearities and turnaround

tff ∼ H−1

Nonlinear regime: self-gravitating gas decoupled from Hubble flow

δρ ρ ∼ a δρ ρ

• δρ

ρ ∼ 1

tff ≡ ✓ 1 16πGρ ◆1/2

Summary of linear perturbation growth

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Primordial regions of under- and over-densities Density contrast grows as Nonlinearities and turnaround

tff ∼ H−1

Nonlinear regime: self-gravitating gas decoupled from Hubble flow

δρ ρ ∼ a δρ ρ

• δρ

ρ ∼ 1

Summary of linear perturbation growth

tff ≡ ✓ 1 16πGρ ◆1/2

STRUCTURE FORMATION HAS TWO MAIN STAGES

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• 1. Linear growth of matter overdensities
• 2. Non-linear evolution of

the resulting(dark) matter clumps

δ = δρ ρ 1

ASTRONOMY BEFORE BIG COMPUTERS

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Fragmentation: Jeans mass decreases as mother halo collapses Jeans Mass: max mass of gas that pressure can support

Low, Linden-Bell 1976 Rees, Ostriker, 1977 Silk, 1977

Halo fragmentation is the origin of stars

M > mJ = π 6 c3

s

✓ π ρ0(z)G ◆3/2 ρeD,

for collapse to happen

JEANS MASS EVOLUTION IS FIXED BY ENERGY CONSERVATION First law of thermodynamics

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dEthermal = −PdV − Λcooling dt

d log TeD d log ρeD = 2 3 meDPeD ρeDTeD − 2 tcollapse tcooling

Specifies a contour in the density-temperature plane as the galaxy collapses

tcooling ≡ 3TeD m 1 Λcooling , tcollapse ≡ ✓d log ρeD dt ◆−1 .

ΛBS = 32α3

DρeDTeD

√πm4

eD

s TeD meD e−mγD /TeD

( )

THE HISTORY OF A GALAXY: SETUP

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We will now follow the evolution of a dark electron halo (1% of our Milky Way)

1010M

THE HISTORY OF A GALAXY: FREE-FALL

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• tff =

✓ 3π 32Gmene ◆1/2

Free-fall time d log T d log ρ = 2 3 Free-fall and heating

THE HISTORY OF A GALAXY: VIRIALIZATION

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• Virialization

d log T d log ρ = 1 3 Cooling time

tcooling ∼ m2

e

α3

Dne

rme Te emγD /Te

THE HISTORY OF A GALAXY: VIRIALIZATION

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• Virialization

d log T d log ρ = 1 3 Cooling time

tcooling ∼ m2

e

α3

Dne

rme Te emγD /Te

THE HISTORY OF A GALAXY: FRAGMENTATION

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• d log T

d log ρ = 2 3 − 2 tff tcooling

Fragmentation Cooling time

tcooling ∼ m2

e

α3

Dne

rme Te emγD /Te

END OF FRAGMENTATION: THE FIRST DARK STARS

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• ★
• *

“Star” here
 means dark-electron
 gas supported by
 repulsive force

Mdark star = 102M neD = 1018cm3 ∼ 106n

A SECOND REASON WHY FRAGMENTATION ENDS

In our model there is a second possibility: fragmentation limited by the dark-photon force

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• ★

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The Sun αD = 1/10

FROM THE LAGRANGIAN TO ASTROPHYSICAL PROPERTIES

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• The dark-electron/photon masses and fine structure constant

set the size of the typical “protostars”

T eD

shock ' 3 ⇥ 103(1 + zta)

 meD 1 MeV  M eD

halo

1010 M 2/3 eV

Bertschinger, ApJS 58, 1985, 39

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Some final remarks

IT IS IMPORTANT TO STUDY BSM STRUCTURE FORMATION

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• []

[⊙]

[ ]

< <

• ≤ ϕ [] ≤ -

α = -

• []

ϕ []

[ ] α = -

≤ 〈σ〉/ ≤ ≤ 〈σ〉/ ≤ 〈 σ

• 〉

/

• =
• 〈

σ

• 〉

/

• =
• Figure 7: Left panel. LIGO best sensitivity (region shaded in green, defined according to
• fig. 2 with DL = 450 Mpc (dashed contour) and DL = 100 Mpc (solid contour)) in terms
• f fermion star mass M and dark matter mass mF . We restrict the analysis to mediator

masses in the range mφ = [10−2 10−1] GeV. The red region is excluded by the condition M > M , which we impose at each point in our sampling of the parameter space. The

• ◆

◆ ◆ ◆ ◆ ◆

R < RS

mχ = 10 GeV mχ = 100 GeV mχ = 1 TeV

10-3 10-2 0.1 1 10-3 10-2 0.1 1 R [km] M [M☉] ρ [ /

3]

• FIG. 3: In the left panels we show dark star mass vs radius relations with DM mass mχ = 10 GeV (Green), 100 GeV

(blue), 1 TeV (purple). Upper, middle and bottom panels correspond to repulsive, no-interactions and attractive interactions respectively. We have fixed µ = 10 MeV and α = 10−3. Solid curves represent full relativistic solutions

Usually, studying only the stability of the ECOs leads to unrealistic setups

1507.00959 1605.01209

SOME THINGS WE CANNOT CALCULATE EASILY

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Minimal mass of ECO Maximal mass of ECO Include accretion (not done here) Compactness Abundance Estimate only: Initial mass function Shape of galaxy/spatial distribution of ECOs

NECO = Mdark galaxy/MECO

INCLUDING RADIATION PRESSURE

Note that our “stars” are all charged! Can their formation

• vercome radiation pressure?

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• ★
• Including radiation pressure

does not significantly modify the results

ΛBS ≥ Ledd M = 4πGmeD σC

PHENOMENOLOGY OF ECOS AND DDM

The phenomenology is similar to PBH phenomenology.

• Microlensing (e.g. Eros 06077207, Kepler 1307.5798)
• Dynamical heating of stellar clusters (Brandt, 1605.03665)
• Dynamical friction (Carr, Sakellariadoi, Apj 516,1999, 195)
• Pulsar timing (Dror et. al. 1901.04490), astrometric lensing (Van Tilburg et. al.1804.01991), fast radio bursts (Muñoz
• et. al. 1605.00008), binaries at GW detectors (Giudice et. al.1605.01209), GW lensing (Jung, Sub Shin1712.01396).

Accretion into baryonic objects (Fan et. Al 1312.1336, Cumberbatch 1005.5102). Is the formation of baryonic structure and dark electron structure correlated? Dark photons could mix with the SM photon. This leads to faint ECOs, can we see them? (Curtin et.al.,

1909.04071/2)

Super-massive BH formation? (D’amico et. Al 1707.03419, Outmezguine et. al. 1807.04750)

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CONCLUSIONS

We described the complete history of structure formation of a simple (the simplest?) dissipative dark sector model. We provided a map between astronomical properties and particle physics parameters. A wide range of opportunities lies ahead,

• What is the behavior of more complicated dark-matter models

with cooling?

• What are the astronomical signatures of such models?
• Numerical simulations?

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Lots of progress to make from the theory side, even if DM interacts with us only gravitationally

AN EXAMPLE UV COMPLETION

Introduce lighter dark proton Oscillations equilibrate the abundance if These components annihilate if

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−yLφ pDpD − yRφ† ¯ pD ¯ pD − mpDpD ¯ pD + h.c. 2yvD = ωosc > H

p+

D, p− D

MD ≥ 10−8 TpD TSM 10−2 αD −3/2mpD keV 210−2 f meD GeV

• eV