sign in sign up
<<
Home , Axino, EDELWEISS

Dark : A Cosmological Perspective

Katie Mack North Carolina State University

www.astrokatie.com : @AstroKatie entire of Physics What We Don’t Know

Origin / particle type Particle mass Thermal history

Non-trivial evolution? One component or many? Non-gravitational interactions (self or SM)? Small-scale behavior (mass of smallest halos) Candidates (incomplete list)

✦ Weakly Interacting Massive (WIMPs)

‣ Something not included in the Standard Model of , generally with weak interactions

‣ May be thermally produced (or not) Annihilating (e.g., SUSY WIMP) Decaying (e.g., sterile ) Warm (WDM) (e.g., ) Self-interacting (SIDM) (particle + dark sector force) (e.g., QCD axion / string axion) Fuzzy DM (tiny mass, large deBroglie wavelength) MACHO (e.g., primordial black holes) WIMP Miracle

Standard thermal WIMP • freezes out when no longer in thermal equilibrium with • for weak-scale mass and cross-section, predict correct abundance of DM • discovery opportunities: annihilation, scattering, production Kolb & Turner 1990 WIMP Direct Detection

10−38 CRESST-II ] DAMA

2 (2015) 5) −40 CDMSLite (201 10 (2017) -2L ICO S P u p e r −42 C 15) D (20 10 M 60 CO- S CoGeNT (2014) PI ) E 015 DE (2 LW -50 (20 E ide ) 16 IS kS 016 ) S Dar (2 −44 ndaX 10 Pa 16) (20 ) 7Be LUX 2017 pp 1T ( ON XEN −46 8 10 B LZ year ton- 200 hep WIN DAR 10−48 Solar-ν

−50 Atm-ν

SI WIMP-nucleon10 cross section [cm SN-ν

10−1 100 101 102 103 104 WIMP mass [GeV/c2] plot via Ciaran O’Hare XENON1T Excess

• may be background, or… • solar / ALPs? (but stellar constraints) • hidden DM? XENON1T Collaboration 2020 • fast (subdominant) DM component? • …? TBD. Directional Detection8

Fluorine recoils [8–50 keVr] September 6 +90 ‣ unique opportunity 5.5 d

b +60

R/ to probe below 4.5 d GeV/c 9

+30 r oa Solar “neutrino floor” Galactic 3.5 [ton 0 plane

Ecliptic 2.5 1 2 30 sr WIMP 1

1.5 ] Galactic latitude, ‣ CYGNUS feasibility 60 0.5 90 paper: Vahsen+2020 180 120 60 0 300 240 180 5 Galactic longitude, l arxiv:2008.12587 FIG. 5. Total angular distribution across the sky of WIMP-induced (blue contours) and neutrino-induced (red contours) fluorine 37 Cygnus 6 yrs recoils on September10 6. We show the distributions in galactic coordinates (l, b) where the line for b =0corresponds to the CRESST ⇥ galactic plane.] 38 Both distributions have been integrated over recoil energy between 8 and 50 keVr.WechooseaWIMPmassof 2 2 10 45 2 8 9GeV c and a SI cross section of 5 10 cm so that its signal is of a similar size to that of the B neutrinos. For reference 39 CDMSlite ⇥ DAMA we also show10 the ecliptic in red: the path along which the Sun, shown by a star, moves over the year. Towards the center of the WIMP recoil distribution we also show the stars of the Cygnus constellation. The blue line which encircles the star Deneb 40 shows the variation10 in the peak direction of the DM wind over the year. PICO2L CYGNUS-KM DarkSide CYGNUS-10 Kamioka, Japan 41 Boulby, UK 10 COSINE-100 CYGNUS 42 3 angular resolution10 of the detector). The separation is unknown) DMPICO60 cross section ,itsprecisevalueisoflow large enough that43 it is possible to discriminateEDELWEISS the two importance before a detection is made. 10 signals even if the recoil vectors cannot be oriented in On the otherPandaX hand the great deal of uncertainty sur- galactic coordinates44 as shown here. This would be the rounding the velocity distribution of the DM is generally 10 He Xenon1T case for experiments45 in which recoil vectors were only more important to understand. The benchmark model 10 F 3 measured after beingXe projected on to a plane [36], or if assumed since the very first direct detection experimentsYGNUS CYGNO 46 10k m C -HD10 timing information10 was unavailable [37]. were conducted is the Standard Halo ModelLead, (SHM), South Dakota in Gran Sasso, Italy 47 which the Milky Way’s3 DM forms an isothermal sphere. 10 The motivation is that it is the simplest model that gives CYGNUS-OZ 48 100k m 10 rise to flat rotation curves. The SHM has a Gaussian dis- Stawell, Aus. tribution for f(v) (or Maxwellian distribution for f(v)), 49 Single threshold: 0.25 keVr [755:5 Torr He:SF6] SI WIMP-proton10 cross section [cm CYGNUS-Andes Worst-case threshold: 8 keVr [755:5 Torr He:SF6] 2 50 3. WIMP astrophysics 1 v Chile/Argentina 10 Search mode: 8 keVr [760 Torr SF6] f(v)= 2 3/2 exp | |2 ⇥(vesc v ) . (9) 51 (2⇡v) Nesc 2v | | 10 1 0 1 2 3 4 ✓ ◆ 10 10 10 10 10 10 Predicting WIMP signals requires astrophysical input Under2 the isothermal SHM, the dispersion is related to in the form of the DM velocity distribution,WIMPf( massv),aswell [GeV/cthe] local rotation speed of circular orbits: v = v0/p2. as the local density of DM ⇢0.Whiletheformerisnot The benchmark used for this speed is the now similarly 1 known concretely, the latter can be determined with as- out-of-date value of v0 220 km s .Amorerecent ⇠ 1 tronomicalFIG. data. 3. Constraints Decades of attempts on the to spin-independent constrain this analysis WIMP- indicates a value (left) of v0 = and 235 km spin-dependent s [82]. As is WIMP- (right) cross sections. 3 value have begun to settle towards ⇢ 0.5 GeV cm convention, the velocity distribution has been truncated We show the existing constraints⇡ and detections from various experiments as labeled (see text for the associated references). In (see,purplee.g. Refs. solid [74–78 and]andRef.[ dashed79]forareviewonmeth- lines we show our projectedat the escape 90% speed CLvesc exclusion, with the limits constant forNesc theusedCygnus to experiment operating for six years ods). However the3 standard value used3 by experimental renormalize the distribution after this truncation.3 Ex- 3 collaborationswith 10 is m0.3upGeV to cm 100,000 .Upcominganalyseswith m of He:SF6 gasperimental at 755:5 Torr analyses (where typically 6 years have assumed1000 mvesc corresponds= 544 to a 1ton-yearexposure).For 1 ⇥ ⇠ the extremelyeach volume high and we precise show statistics limits offor the two astromet- possibleor thresholds,533 km s ,thelatterbeingthemorerecentRAVE ranging from the worst-case electron discrimination threshold of 8 keVr ric survey Gaia [80]areexpectedtoleadtoevenmore to a very best-case minimum threshold corresponding to a single electron, 0.25 keVr.Weemphasizehoweverthatweanticipate tightly constrained values in the next few years. Since 3 the localelectron dark matterdiscrimination density only well amounts below to 8 a keV mul-r.Forthe100km limits we add a third dotted line which corresponds to a mode with 3 tiplicativepurely factor SF6 whichgas can at 760 be absorbed Torr. This into the experiment (also Although, would see have Ref. [81 a] significantlyfor a specific case where higher this is total not true. mass but would come at the cost of any directional sensitivity. This ‘search mode’ could be used to extend the high mass sensitivity to just within reach of the neutrino floor. For the SI panel, we shade in gray the neutrino floor for helium, fluorine, and targets (top to bottom), and for SD we show only fluorine and xenon. We define the neutrino floor as the cross section limit at which the rate of improvement with increasing exposure is the slowest in standard direct detection–the effect that Cygnus aims to circumvent. This definition corresponds to (100) neutrino events. O

WIMP velocity, v, with the outgoing recoil direction ˆr, be such a powerful means to discover DM. The pri- mary signal is a dipole anisotropy towards the direction m E ˆr = vˆ ,leadingtoan (10) forward-backward asym- ˆ = A r v , lab O v r 2 min (5) metry in the number of recoil events. The strength of · s 2µA ⌘ this dipole means that in ideal circumstances (i.e. perfect which can be enforced in the differential cross section recoil direction reconstruction) an isotropic null hypoth- with a delta function, esis for the recoil direction distribution can be rejected at 90% confidence with only (10) events [43, 50]. With 2 O d d 1 (30) recoil directions it becomes possible to point back = v (v ˆr vmin) , (6) O dEr d⌦ dEr 2⇡ · towards Cygnus and confirm the galactic origin of the sig- nal [44]. Secondary signals such as a ring feature at low where d⌦ is the solid angle element around ˆr. The event energies [51], and the aberration of recoil directions over rate then has the structure, time [52], may also aid in the confirmation of a DM sig- 2 nal, as well as for characterizing astrophysical properties d R ⇢0 SI 2 SD 2 (Er, ˆr,t)= 2 (0 FSI(E)+0 FSD(E)) of the DM halo. dErd ⌦ 4⇡µpm

fˆ(vmin, ˆr,t) , (7) ⇥ Existing limits and projections for Cygnus where the velocity distribution now enters in the form of its Radon transform [48, 49], Figure 3 shows a selection of constraints from di- rect detection experiments along with our headline re- ˆ 3 f(vmin, ˆr,t)= (v ˆr vmin) f(v,t) d v . (8) sult: the WIMP reach of Cygnus. Constraints ex- · 2 Z ist for WIMPs with masses larger than 1 GeV c 46⇠ 2 The characteristic angular structure of the DM flux and SI cross sections larger than 10 cm .Under- ⇠ on Earth is the reason why directional detection could neath these limits lies the neutrino floor, below which Candidates (incomplete list)

✦ Weakly Interacting Massive Particles (WIMPs)

‣ Something not included in the Standard Model of Particle Physics, generally with weak interactions

‣ May be thermally produced (or not) Annihilating (e.g., SUSY neutralino WIMP) Decaying (e.g., ) Warm (WDM) (e.g., axino) Self-interacting (SIDM) (particle + dark sector force)

Axion (e.g., QCD axion / string axion) (Mack 2011; Mack & Steinhardt 2011) Fuzzy DM (tiny mass, large deBroglie wavelength)

MACHO (e.g., primordial black holes) (Mack, Ostriker & Ricotti 2007; R,O,M 2008) Candidates (incomplete list)

6 M.R. Buckley, A.H.G. Peter / Physics Reports 761 (2018) 1–60 halo mass

1 Fig. 1. Estimates for the range of particle physics and astrophysics figures of merit (⇤ and Mhalo) for a variety of dark matter models. The range of Mhalo covered by ‘‘evolutionary’’ and ‘‘primordial’’ self-interacting dark matter models (SIDM) are overlapping. The former covers the range 106–1015 M , and the latter the range below 1011 M . SeeStandard text for further details. Model interaction

these(where questions, we the expect answer to is see a qualified a deviation yes. For from example, CDM) stars in the visible sector might actPeter as microlenses & Buckley for dark 2018 matter astronomers, as they do for visible-sector applications [53]. The abundance and stability of such objects may hint at nuclear fusion-type processes. But for some of these questions – e.g., three generations of and – the answer might be no. This thought experiment is not an attempt to argue that a dark-matter scientist could immediately determine the unique properties of the baryons. As we demonstrated, there are several possible branches that our hypothetical researcher may wander down; it is not clear that all other options would not also lead to self-consistent results. Indeed, one might expect the same level of vigorous debate as to the nature of this mysterious extra component of the Universe as found among our baryonic theoretical physicists. However, the dark-matter scientist would be able to map out some of the most important features of the Standard Model, like electromagnetism, which were also the first Standard Model features that were described by modern theory by visible-sector scientists.

3. Metrics for dark matter models

As our thought experiment demonstrates, much may be learned about the complicated Standard Model particle physics through measurements of the gravitational imprint of baryons if we were dark-matter scientists surveying the Universe. We can uncover non-trivial dark matter physics in the same manner. A comprehensive characterization of dark matter microphysics requires a combination of approaches: laboratory-based particle physics searches for interactions with the Standard Model, and the astronomical searches for interactions within a dark sector and also (as we will see) with the Standard Model. To organize these searches, we need a compact space in which to classify models in terms of their observability in the laboratory and in the sky. Our goal with this section is to motivate a specific choice for this space, and to show how particle dark matter models inhabit it. The space is designed to be well-matched to the ways particle physicists and astronomers think about dark matter, making the mapping between the particle and astronomical spaces transparent and straight-forward, and compact but informative enough so that one might define ‘‘figures of merit’’ to quantify how well future experiments and observations will constrain dark matter models. 1 We classify dark matter models by their interaction strength with the Standard Model, ⇤ , and the cosmological scale at which we expect to see a deviation from the (CDM) paradigm, Mhalo. The former defines the sensitivity of particle physics detectors and the latter defines the largest size of the systems that must be understood in order to discover model-specific dark matter structures. We consider each axis of this parameter space in turn, and classify some well-known dark matter models by where they fall in the resulting two-dimensional parameter space. We summarize our estimates for these models in Fig. 1, with details given in the text. Because one of our goals is to enable better communication between Possible Hints/Signals Annihilation?

Gamma rays in the Excess at Excess 3 Galactic Center high energy III. ANTIPROTONat high AND energy FITTINGS

FIG. 2: (a) Positron fraction (solid line), which includes the and positrons coming from the DC and back- ground electrons (dotted line, for example see Refs. [29, 30]). Filled circles correspond to the AMS-02 data [1, 34, 35] and FIG. 1: fraction fitted to the data. The data PAMELA data [5] (b) Total electron and positron flux (solid AMS Collaboration 2013 points are taken byKohri [1] for AMS-02,et al. and2015 by [15] for PAMELA. line). The flux of the electrons and positrons created only in Daylan et al. 2014 The dotted line is plotted only by using the background the DC (background) is plotted by the dashed (dotted) line. flux [33]. The shadow region represents the uncertainties of Observational data by AMS-02, Fermi, HESS, BETS, PPB- the background flux among the propagation models shown BETS, and ATIC2 [6–8, 36] are also plotted. The shadow in [1]. Not pulsars! region represents the uncertainty of the HESS data. … but maybe pulsars In Fig. 1, we plot the antiproton fraction at the Earth fitted only by adding astrophysical local contributions in our model (See the model B shown in Ref. [4]). For produced from the same pp collision sources. the background… flux, webut adopted maybe the 15% smaller value of the mean value shown in [33]. Here, the radius of a sphericalsupernova DC, RDC = 40 pc is adopted. remnant The targetproton −3 IV. CONCLUSION density is set to be n0 =50cm . The spectral index s =1.75 and the maximum energy Emax =100TeVare assumed. We take the duration of the pp collision to be We have discussed the anomaly of the antiproton frac- 5 tpp =2× 10 yr. The total energy of the accelerated tion recently-reported by the AMS-02 experiment. By 50 is assumed to be Etot,p =3× 10 erg. The considering the same origin of the pp collisions between distance to the front of the DC is set to be d =200pc. cosmic-ray protons accelerated by SNRs and a dense − + 5 About the diffusion time of e and e , tdiff =2× 10 yr cloud which surrounds the SNRs, we can fit the data is adopted. We take the magnetic field outside the DC of the observed antiproton and positron simultaneously to be Bdiff =3µG (See [4] for the further details). without a fine tuning in the model parameters. The ob- In Fig. 2 we also plot the positron fraction and the total served fluxes of both antiprotons and positrons are con- e−+e+ flux. It is remarkable that we can automatically sistent with our predictions shown in Ref. [4]. fit the observational data of both the positron fraction Regardless of the model details, the ratio of antipro- and the total e− + e+ flux by using the same set of the tons and positrons is essentially determined by the fun- parameters [4]. damental branching ratio of the pp collisions. Thus the The positron fraction rises at higher energies than that observed antiproton excess should entail the positron ex- of the antiproton fraction (Fig. 2), because the spectral cess, and vice versa. This does not depend on the propa- index of the background antiproton is harder than that of gation model since both propagate in a sim- the background positron. This comes from a difference ilar way below the cooling cutoff energy ∼ TeV. between their cooling processes. Only for background The cutoff energy of e− cooling marks the supernova positrons and electrons the cooling is effective in the cur- age of ∼ 105 years [18, 37], while we also expect a e+ rent situation. cutoff. The trans-TeV energy will be probed by the fu- In Fig. 3, we plot the positron to antiproton ratio as a ture CALET, DAMPE and CTA experiments [38, 40]. function of the rigidity. Here the local components repre- An anisotropy of the arrival direction is also a unique sent the contribution of the nearby SNRs produced only signature, e.g., [39]. by the pp collisions. From this figure, we find that both The boron to carbon ratio as well as the Li to car- of the positron and the antiproton can be consistently bon ratio have no clear excesses [1]. This suggests that Decay?

Excess x-rays in galaxy clusters

Bulbul et al. 2014

… but maybe line contamination Scattering?

Super-cold neutral hydrogen at high redshift x-rays

stars

Bowman et al. 2018 Pritchard & Loeb 2010 … but maybe a foreground subtraction problem The Cosmic Frontier Dark Matter: Cosmology

Paul Angel, Tiamat Simulation Impact of Dark Matter Annihilation

Major unanswered question: If dark matter annihilates across all of cosmic time, how does it affect the first stars and galaxies?

heating

gamma rays Lyman(alpha/Werner) Annihilation in the Intergalactic Medium

halo halo

halo halo Annihilation in the Intergalactic Medium

halo halo

halo halo Usual treatment: • monolithic halos • immediate uniform energy deposition Annihilation in the Intergalactic Medium

inverse Compton scattering

Better: • structured halos • delayed energy deposition Annihilation Feedback on Halo Gas

If dark matter is annihilating within baryonic halos, does this constitute an effective “feedback” process?

PYTHIA code: dark matter annihilation events MEDEA2 code: energy transfer to baryons Halo models: density profile, mass-concentration Annihilation Feedback on Halo Gas

50 Comparing: 40 dark matter

log annihilation

1 10 energy

30 ( (over Hubble time) F

eff to:

0 )

redshift (z) redshift 20 gas binding -1 -3 energy 10 -2 103 105 107 109 mass (MSun)

Schon, Mack+ 2015, MNRAS [arxiv: 1411.3783] Annihilation Feedback on Halo Gas

50 Comparing: 40 dark matter

stars log annihilation

1 form 10 energy

30 ( (over Hubble time) here F eff to:

0 )

redshift (z) redshift 20 gas binding -1 -3 energy 10 -2 103 105 107 109 mass (MSun)

Schon, Mack+ 2015, MNRAS [arxiv: 1411.3783] Annihilation Feedback on Halo Gas

50 Comparing: dark matter 40 2

stars log annihilation

form 10 energy

30 ( 1 (over Hubble time) here F eff to: )

redshift (z) redshift 20 gas binding 0 energy 10 -1

103 105 107 109 mass (MSun) Schon, Mack+ 2015, MNRAS [arxiv: 1411.3783] + Schon, Mack+ 2018, MNRAS [arxiv: 1706.04327] Probing Cosmic Dawn

Djorgovski et al., Caltech current instruments next decade

JWST

SKA Dark Matter & 21cm

List, Elahi and Lewis 2020

DM annihilation during Cosmic Dawn z=11 5

Evoli et al. 2014 Annihilating dark matter can heat and ionize the IGM, Figure 3. Gas temperature (top) and 21cm brightness temperature (bottom )atz = 11. The light blue spheres are located at altering the 21cm signal at the centers of mass of dark matter haloes and their radii correspond to 5 Rvir.Form = 10 GeV, the cold HI gas is universally heated, causing the 21cm line to be in emission everywhere. For a m = 100 GeV particle, DM heating is strong enough to induce Tb > 0 in vast parts of the filamentary structure, whereas Tb . 0 in voids. Without DM annihilation, only gas in cosmic dawn dense regions (and in particular in haloes) can be seen in emission. An animated version of this figure is available in the HTML version of this paper,DM+Hydro which shows the simulation boxes simulations from di↵erent angles and zoomsneeded into the boxes. to

ature in the entire simulation box, and the 21cm line a minimum temperature floor of Tgas = 10 K is set for is globally in emission.trace For the the 100 GeV impact WIMP, only ofnumerical DM stability, annihilation for which reason the lower end of the (and even dominate heating at voids remain in absorption, while large regions are in Tb distribution should not be over-interpreted. emission. In contrast,on the galactic volume-averaged brightnessand intergalactic gas temperature in the case without DM annihilation is neg- 4. DISCUSSION certain redshifts) ative, and only gas within filamentary structures is seen In this paper, we have presented first results from a in emission. suite of hydrodynamic simulations that self-consistently This is confirmed by considering the distribution of include annihilating DM in conjunction with baryonic brightness temperature values per cell, depicted in Fig- cooling physics at high redshift. We have analyzed how ure 4. The distribution is computed by mapping the the spatial distribution and the T ⇢ phase space dis- brightness temperature values from the N-body gas par- tribution of the hydrogen gas at the end of Cosmic ticles onto a regular grid of size 5123 using Shepard in- Dawn are a↵ected by annihilating DM. Running hydro- terpolation, binning the resulting values, and using a dynamic simulations instead of resorting to approximate kernel density estimator. The distribution of T peaks b methods allows us to evaluate the power generated by at 24, 3, and 13 mK without DM annihilation, for ⇠ DM annihilation as a function of the non-linear local m = 100 GeV, and for m = 10 GeV, respectively, DM density field and to locally deposit the result DM while the means are located at 6, 9, and 18 mK. ⇠ heating, without the need for analytic halo mass func- The brightness temperature distribution becomes nar- tions or halo profiles. This is in contrast to previous rower as the DM mass decreases. This is in line with investigations of the DM annihilation imprint on the findings by Vald´es et al. (2013); Evoli et al. (2014)who 21cm brightness temperature in literature that assume show that the 21cm power spectrum is reduced due to a redshift-dependent but spatially homogeneous boost the relative uniformity of DM heating as compared to factor calculated by integrating over analytic models for heating from astrophysical sources. In our simulations, the and halo profiles, leading to spa- dark matter annihilates

internal heating radiation of dark matter background halos

low-mass cut-off

structure of first stars heating/ ionization of intergalactic gas

H2 abundance black hole formation JWST unified simulation of galaxy formation & evolution

high-redshift high-redshift 21cm signal galaxies Image credit: Swinburne/ credit: Image ICRAR/Cambridge/ASTRON SKA

Image credit: NASA Take-Home Messages

Future surveys can probe the particle physics of dark matter and produce a more consistent picture of cosmology To determine dark matter’s impact on high-redshift astrophysics, we need to understand small halos and their evolution end