1

25. Cosmological Parameters

25. Cosmological Parameters
Revised August 2023 by O. Lahav (UCL) and A.R. Liddle (ULisboa).

25.1 Parametrizing the Universe
Rapid advances in observational cosmology have led to the establishment of a precision cosmological model, with many of the key cosmological parameters determined to one or two significant
figure accuracy. Particularly prominent are measurements of cosmic microwave background (CMB)
anisotropies, with the highest precision observations being those of the Planck Satellite [1,2], which
supersede the landmark WMAP results [3, 4]. However the most accurate model of the Universe
requires consideration of a range of observations, with complementary probes providing consistency
checks, lifting parameter degeneracies, and enabling the strongest constraints to be placed.
The term ‘cosmological parameters’ now has a wide scope, and may include the parameterizationof some functions as well as simple numbers describing properties of the Universe. The original usagereferred to the parameters describing the global dynamics of the Universe, such as its expansion
rate and curvature. Now we wish to know how the matter budget of the Universe is built up
from its constituents: baryons, photons, neutrinos, dark matter, and dark energy. We also needto describe the nature of perturbations in the Universe, through global statistical descriptors suchas the matter and radiation power spectra. There may be additional parameters describing thephysical state of the Universe, such as the ionization fraction as a function of time during theera since recombination. Typical comparisons of cosmological models with observational data now
feature between five and ten parameters.
25.1.1 The global description of the Universe
Ordinarily, the Universe is taken to be a perturbed Robertson–Walker space-time, with dynamics governed by Einstein’s equations. This is described in detail in the Big-Bang Cosmology
chapter in this volume. Using the density parameters Ωi for the various matter species and ΩΛ forthe cosmological constant, the Friedmann equation can be writtenXiΩi + ΩΛ − 1 =k
R2 H 2,(25.1)

where the sum is over all the different species of material in the Universe. This equation applies atany epoch, but later in this article we will use the symbols Ωi and ΩΛ to refer specifically to thepresent-epoch values.
The complete present-epoch state of the homogeneous Universe can be described by giving thecurrent-epoch values of all the density parameters and the Hubble constant h (the present-day
Hubble parameter being written H0 = 100h km s−1 Mpc−1 ). A typical collection would be baryonsΩb , photons Ωγ , neutrinos Ων , and cold dark matter Ωc (given charge neutrality, the electron
density is guaranteed to be too small to be worth considering separately and is effectively included
with the baryons). The spatial curvature can then be determined from the other parameters using
Eq. (25.1). The total present matter density Ωm = Ωc + Ωb may be used in place of the cold dark
matter density Ωc .
These parameters also allow us to track the history of the Universe, at least back until an epoch
where interactions allow interchanges between the densities of the different species; this is believed
to have last happened at neutrino decoupling, shortly before Big-Bang Nucleosynthesis (BBN). To
probe further back into the Universe’s history requires assumptions about particle interactions, andperhaps about the nature of physical laws themselves.
S. Navas et al. (Particle Data Group), Phys. Rev. D 110, 030001 (2024)
31st May, 2024 10:15am

2

25. Cosmological Parameters

The standard neutrino sector has three flavors. For neutrinos of mass in the range 5 × 10−4 eV
to 1 MeV, the density parameter in neutrinos is predicted to beΩν h2 =

Pmν,93.12 eV(25.2)

where the sum is over all families with mass in that range (higher masses need a more sophisticatedcalculation). We use units with c = 1 throughout. Results on atmospheric and Solar neutrino
oscillations [5] imply non-zero mass-squared differences between the three neutrino flavors. Theseoscillation experiments cannot tell us the absolute neutrino masses, but within the normal assumption of a mass hierarchy suggest a lower limit of approximately 0.06 eV for the sum of the neutrino
masses (see the Neutrino chapter).
Even a mass this small has a potentially observable effect on the formation of structure, as
neutrino free-streaming damps the growth of perturbations. Analyses commonly now either assume
a neutrino mass sum fixed at this lower limit, or allow the neutrino mass sum to be a variable
parameter. To date there is no decisive evidence of any effects from either neutrino masses oran otherwise non-standard neutrino sector, and observations impose quite stringent limits; see the
Neutrinos in Cosmology chapter. However, we note that the inclusion of the neutrino mass sum as
a free parameter can affect the derived values of other cosmological parameters.
25.1.2 Inflation and perturbations
A complete model of the Universe must include a description of deviations from homogeneity,
at least statistically. Indeed, the most powerful probes of the parameters described above comefrom the evolution of perturbations, so their study is naturally intertwined with the determinationof cosmological parameters.
There are many different notations used to describe the perturbations, both in terms of thequantity used and the definition of the statistical measure. We use the dimensionless power spectrum ∆2 as defined in the Big Bang Cosmology section (also denoted P in some of the literature).
If the perturbations obey Gaussian statistics, the power spectrum provides a complete descriptionof their properties.
From a theoretical perspective, a useful quantity to describe the perturbations is the curvature
perturbation R, which measures the spatial curvature of a comoving slicing of the space-time.
A simple case is the Harrison–Zeldovich spectrum, which corresponds to a constant ∆2R . Moregenerally, one can approximate the spectrum by a power law, writing∆2R (k) = ∆2R (k∗ )k ns −1,k∗(25.3)

where ns is known as the spectral index, always defined so that ns = 1 for the Harrison–Zeldovich
spectrum, and k∗ is an arbitrarily chosen scale. The initial spectrum, defined at some early epochof the Universe’s history, is usually taken to have a simple form such as this power law, and we will
see that observations require ns close to one. Subsequent evolution will modify the spectrum fromits initial form.
The simplest mechanism for generating the observed perturbations is the inflationary cosmology,
which posits a period of accelerated expansion in the Universe’s early stages [6, 7]. It is a usefulworking hypothesis that this is the sole mechanism for generating perturbations, and it may further
be assumed to be the simplest class of inflationary model, where the dynamics are equivalent tothat of a single scalar field φ with canonical kinetic energy slowly rolling on a potential V (φ). One
may seek to verify that this simple picture can match observations and to determine the properties
of V (φ) from the observational data. Alternatively, more complicated models, perhaps motivated31st May, 2024

25. Cosmological Parameters

3by contemporary fundamental physics ideas, may be tested on a model-by-model basis (see morein the Inflation chapter in this volume).
Inflation generates perturbations through the amplification of quantum fluctuations, which
are stretched to astrophysical scales by the rapid expansion. The simplest models generate two
types, density perturbations that come from fluctuations in the scalar field and its corresponding
scalar metric perturbation, and gravitational waves that are tensor metric fluctuations. The former
experience gravitational instability and lead to structure formation, while both generate CMB
anisotropies. Defining slow-rollparameters (with primes indicating derivatives with respect to thepscalar field, and mPl ≡ ~c/G the Planck mass) asm2
 = Pl
16π



V0
V

2

,

η=

m2Pl V 00,
8π V(25.4)

which should satisfy , |η|  1, the spectra can be computed using the slow-roll approximation as∆2R (k) '8 V
3m4Pl  k=aH,

∆2t (k) '128.
V
3m4Pl k=aH

(25.5)

In each case, the expressions on the right-hand side are to be evaluated when the scale k is equal to
the Hubble radius during inflation. The symbol ‘'’ here indicates use of the slow-roll approximation,
which is expected to be accurate to a few percent or better.
From these expressions, we can compute the spectral indices [8]:ns ' 1 − 6 + 2η;nt ' −2 .

(25.6)

Another useful quantity is the ratio of the two spectra, defined byr≡∆2t (k∗ ).
∆2R (k∗ )(25.7)We haver ' 16 ' −8nt ,(25.8)

which is known as the consistency equation.
One could consider corrections to the power-law approximation, which we discuss later. However, for now we make the working assumption that the spectra can be approximated by such powerlaws. The consistency equation shows that r and nt are not independent parameters, and so thesimplest inflation models give initial conditions described by three parameters, usually taken as∆2R , ns , and r, all to be evaluated at some scale k∗ , usually the ‘statistical center’ of the range
explored by the data. Alternatively, one could use the parametrization V , , and η, all evaluatedat a point on the putative inflationary potential.
After the perturbations are created in the early Universe, they undergo a complex evolution
up until the time they are observed in the present Universe. When the perturbations are small,this can be accurately followed using a linear theory numerical code such as CAMB or CLASS [9].
This works right up to the present for the CMB, but for density perturbations on small scales nonlinear evolution is important and can be addressed by a variety of semi-analytical and numerical
techniques. However the analysis is made, the outcome of the evolution is in principle determinedby the cosmological model and by the parameters describing the initial perturbations, and hence
can be used to determine them.
Of particular interest are CMB anisotropies. Both the total intensity and two independent
polarization modes are predicted to have anisotropies. These can be described by the radiation
angular power spectra C` as defined in the CMB article in this volume, and again provide a complete
description if the density perturbations are Gaussian.
31st May, 2024

25. Cosmological Parameters425.1.3 The standard cosmological model
We now have most of the ingredients in place to describe the cosmological model. Beyond thoseof the previous subsections, we need a measure of the ionization state of the Universe. The Universe
is known to be highly ionized at redshifts below 5 or so (otherwise radiation from distant quasars
would be heavily absorbed in the ultra-violet), and the ionized electrons can scatter microwave
photons, altering the pattern of observed anisotropies. The most convenient parameter to describe
this is the optical depth to scattering τ (i.e., the probability that a given photon scatters once); in
the approximation of instantaneous and complete reionization, this could equivalently be describedby the redshift of reionization zi .
As described in Sec. 25.4, models based on these parameters are able to give a good fit to
the complete set of high-quality data available at present, and indeed some simplification is possible. Observations are consistent with spatial flatness, and the inflation models so far described
automatically generate negligible spatial curvature, so we can set k = 0; the density parametersthen must sum to unity, and so one of them can be eliminated. The neutrino energy density isoften not taken as an independent parameter; provided that the neutrino sector has the standard
interactions, the neutrino energy density, while relativistic, can be related to the photon density
using thermal physics arguments, and a minimal assumption takes the neutrino mass sum to bethat of the lowest mass solution to the neutrino oscillation constraints, namely 0.06 eV. In addition,
there is no observational evidence for the existence of tensor perturbations (with the upper limits
now starting to become constraining on models), and so r could be set to zero. This leaves sevenparameters, which is the smallest set that can usefully be compared to the present cosmological
data. This model is referred to by various names, including ΛCDM, the concordance cosmology,
and the standard cosmological model.
Of these parameters, only Ωγ is accurately measured directly. The radiation density is dominated by the energy in the CMB, and the COBE satellite FIRAS experiment determined its
temperature to be T = (2.7255 ± 0.0006) K [10],1 corresponding to Ωγ = 2.47 × 10−5 h−2 . It typically can be taken as fixed when fitting other data. Hence the minimum number of cosmologicalparameters varied in fits to data is six, though as described below there may additionally be many‘nuisance’ parameters necessary to describe astrophysical processes influencing the data.
In addition to this minimal set, there is a range of other parameters that might prove importantin future as the datasets further improve, but for which there is so far no direct evidence, allowing
them to be set to specific values for now. We discuss various speculative options in the next
section. For completeness at this point, we mention one other interesting quantity, the helium
fraction, which is a non-zero parameter that can affect the CMB anisotropies at a subtle level. It is
usually fixed in microwave anisotropy studies, but the data are approaching a level where allowing
its variation may become mandatory.
In conventional parameter estimation, a set of parameters is chosen by hand and the aim is
to constrain their values. The higher-level inference problem of model selection instead compares
different choices of parameter sets, as is necessary to assess whether observations are pointing
towards inclusion of new physical effects. Bayesian inference offers an attractive framework forcosmological model selection, setting a tension between model predictiveness and ability to fit thedata [11], and its use is becoming widespread.
25.1.4 Derived parameters
The parameter list of the previous subsection is sufficient to give a complete description of cosmological models that agree with observational data. However, it is not a unique parameterization,
1

All quoted uncertainties in this article are 1σ/68% confidence and all upper limits are 95% confidence. Cosmological parameters sometimes have significantly non-Gaussian uncertainties. Values from original sources have beenrounded according to the conventions of this volume.
31st May, 2024

525. Cosmological Parameters

and one could instead use parameters derived from that basic set. Parameters that can be obtained
from the set given above include the age of the Universe, the present horizon distance, the present
neutrino background temperature, the epoch of matter–radiation equality, the epochs of recombination and decoupling, the epoch of transition to an accelerating Universe, the baryon-to-photon
ratio, and the baryon-to-dark-matter density ratio. In addition, the physical densities of the mattercomponents, Ωi h2 , are often more useful than the density parameters. The density perturbation
amplitude can be specified in many different ways other than the large-scale primordial amplitude,
for instance, in terms of its effect on the CMB, or by specifying a short-scale quantity, a common
choice being the present linear-theory mass dispersion in a radius of 8 h−1 Mpc, known as σ8 .
Different types of observation are sensitive to different subsets of the full cosmological parameter
set, and some are more naturally interpreted in terms of some of the derived parameters of thissubsection than on the original base parameter set. In particular, most types of observation featuredegeneracies whereby they are unable to separate the effects of simultaneously varying specificcombinations of several of the base parameters.

25.2 Extensions to the standard model
At present, there is no positive evidence in favor of extensions of the standard model. These
are becoming increasingly constrained by the data, though there always remains the possibility oftrace effects at a level below present observational capability.
25.2.1 More general perturbations
The standard cosmology assumes adiabatic, Gaussian perturbations. Adiabaticity means that
all types of material in the Universe share a common perturbation, so that if the space-time isfoliated by constant-density hypersurfaces, then all fluids and fields are homogeneous on those slices,with the perturbations completely described by the variation of the spatial curvature of the slices.
Gaussianity means that the initial perturbations obey Gaussian statistics, with the amplitudes
of waves of different wavenumbers being randomly drawn from a Gaussian distribution of width
given by the power spectrum. Note that gravitational instability generates non-Gaussianity; in thiscontext, Gaussianity refers to a property of the initial perturbations, before they evolve.
The simplest inflation models, based on one dynamical field, predict adiabatic perturbations
and a level of non-Gaussianity that is too small to be detected by any experiment so far conceived.
For present data, the primordial spectra are usually assumed to be power laws.
25.2.1.1 Non-power-law spectra
For typical inflation models, it is an approximation to take the spectra as power laws, albeit
usually a good one. As data quality improves, one might expect this approximation to come underpressure, requiring a more accurate description of the initial spectra, particularly for the density
perturbations. In general, one can expand ln ∆2R asln ∆2R (k) = ln ∆2R (k∗ ) + (ns,∗ − 1) ln

k
1 dns
k
+
ln2
+ ··· ,k∗ 2 d ln k ∗k∗(25.9)

where the coefficients are all evaluated at some scale k∗ . The term dns /d ln k|∗ is often called
the running of the spectral index [12]. Once non-power-law spectra are allowed, it is necessary to
specify the scale k∗ at which the spectral index is defined.
25.2.1.2 Isocurvature perturbations
An isocurvature perturbation is one that leaves the total density unperturbed, while perturbing
the relative amounts of different materials. If the Universe contains N fluids, there is one growing
adiabatic mode and N − 1 growing isocurvature modes (for reviews see Ref. [7] and Ref. [13]).
These can be excited, for example, in inflationary models where there are two or more fields that
31st May, 2024

625. Cosmological Parameters

acquire dynamically-important perturbations. If one field decays to form normal matter, while
the second survives to become the dark matter, this will generate a cold dark matter isocurvature
perturbation.
In general, there are also correlations between the different modes, and so the full set of perturbations is described by a matrix giving the spectra and their correlations. Constraining such ageneral construct is challenging, though constraints on individual modes are beginning to become
meaningful, with no evidence that any other than the adiabatic mode must be non-zero.
25.2.1.3 Seeded perturbations
An alternative to laying down perturbations at very early epochs is that they are seeded throughout cosmic history, for instance by topological defects such as cosmic strings. It has long been
excluded that these are the sole original of structure, but they could contribute part of the perturbation signal, current limits being just a few percent [14]. In particular, cosmic defects formed ina phase transition ending inflation is a plausible scenario for such a contribution.
25.2.1.4 Non-Gaussianity
Multi-field inflation models can also generate primordial non-Gaussianity (reviewed, e.g., in
Ref. [7]). The extra fields can either be in the same sector of the underlying theory as the inflaton,or completely separate, an interesting example of the latter being the curvaton model [15]. Current
upper limits on non-Gaussianity are becoming stringent, but there remains strong motivation topush down those limits and perhaps reveal trace non-Gaussianity in the data. If non-Gaussianity isobserved, its nature may favor an inflationary origin, or a different one such as topological defects.
25.2.2 Dark matter properties
Dark matter properties are discussed in the Dark Matter chapter in this volume. The simplest
assumption concerning the dark matter is that it has no significant interactions with other matter,and that its particles have a negligible velocity as far as structure formation is concerned. Suchdark matter is described as ‘cold,’ and candidates include the lightest supersymmetric particle, theaxion, and primordial black holes. As far as astrophysicists are concerned, a complete specificationof the relevant cold dark matter properties is given by the density parameter Ωc , though thoseseeking to detect it directly need also to know its interaction properties.
Cold dark matter is the standard assumption and gives an excellent fit to observations, except
possibly on the shortest scales where there remains some controversy concerning the structure of
dwarf galaxies and possible substructure in galaxy halos. It has long been excluded for all the dark
matter to have a large velocity dispersion, so-called ‘hot’ dark matter, as it does not permit galaxies
to form; for thermal relics the mass must be above about 1 keV to satisfy this constraint, thoughrelics produced non-thermally, such as the axion, need not obey this limit. However, in future
further parameters might need to be introduced to describe dark matter properties relevant toastrophysical observations. Suggestions that have been made include a modest velocity dispersion(warm dark matter) and dark matter self-interactions. There remains the possibility that the dark
matter is comprised of two separate components, e.g., a cold one and a hot one, an example being
if massive neutrinos have a non-negligible effect.
25.2.3 Relativistic species
The number of relativistic species in the young Universe (omitting photons) is denoted Neff . In
the standard cosmological model only the three neutrino species contribute, and its baseline value
is assumed fixed at 3.044 (the small shift from 3 is because of a slight predicted deviation from athermal distribution [16]). However other species could contribute, for example an extra neutrino,
possibly of sterile type, or massless Goldstone bosons or other scalars. It is hence interesting tostudy the effect of allowing this parameter to vary, and indeed although 3.044 is consistent with
31st May, 2024

725. Cosmological Parameters

the data, most analyses currently suggest a somewhat higher value (e.g., Ref. [17]).
25.2.4 Dark energy and modified gravity
While the standard cosmological model given above features a cosmological constant, in order toexplain observations indicating that the Universe is presently accelerating, further possibilities exist
under the general headings of ‘dark energy’ and ‘modified gravity’. These topics are described indetail in the Dark Energy chapter in this volume. This article focuses on the case of the cosmologicalconstant, since this simple model is a good match to existing data. We note that more generaltreatments of dark energy/modified gravity will lead to weaker constraints on other parameters.
25.2.5 Complex ionization history
The full ionization history of the Universe is given by the ionization fraction as a functionof redshift z. The simplest scenario takes the ionization to have the small residual value left
after recombination up to some redshift zi , at which point the Universe instantaneously reionizes
completely. Then there is a one-to-one correspondence between τ and zi (that relation, however,also depending on other cosmological parameters). An accurate treatment of this process will track
separate histories for hydrogen and helium. While currently rapid ionization appears to be a good
approximation, as data improve a more complex ionization history may need to be considered.
25.2.6 Varying ‘constants’
Variation of the fundamental constants of Nature over cosmological times is another possible
enhancement of the standard cosmology. There is a long history of study of variation of thegravitational constant GN , and more recently attention has been drawn to the possibility of small
fractional variations in the fine-structure constant. There is presently no observational evidence forthe former, which is tightly constrained by a variety of measurements. Evidence for the latter has
been claimed from studies of spectral line shifts in quasar spectra at redshift z ' 2 [18], but this ispresently controversial and in need of further observational study.
25.2.7 Cosmic topology
The usual hypothesis is that the Universe has the simplest topology consistent with its geometry,for example that a flat Universe extends forever. Observations cannot tell us whether that is true,
but they can test the possibility of a non-trivial topology on scales up to roughly the present Hubble
scale. Extra parameters would be needed to specify both the type and scale of the topology; forexample, a cuboidal topology would need specification of the three principal axis lengths andorientation. At present, there is no evidence for non-trivial cosmic topology [19].

25.3 Cosmological Probes
The goal of the observational cosmologist is to utilize astronomical information to derive cosmological parameters. The transformation from the observables to the parameters usually involves
many assumptions about the nature of the data, as well as of the dark sector. Below we outline thephysical processes involved in each of the major probes, and the main recent results. The first two
subsections concern probes of the homogeneous Universe, while the remainder consider constraints
from perturbations.
In addition to statistical uncertainties we note three sources of systematic uncertainties that will
apply to the cosmological parameters of interest: (i) due to the assumptions on the cosmologicalmodel and its priors (i.e., the number of assumed cosmological parameters and their allowed range);(ii) due to the uncertainty in the astrophysics of the objects (e.g., light-curve fitting for supernovae
or the mass–temperature relation of galaxy clusters); and (iii) due to instrumental and observational
limitations (e.g., the effect of ‘seeing’ on weak gravitational lensing measurements, or beam shapeon CMB anisotropy measurements).
These systematics, the last two of which appear as ‘nuisance parameters’, pose a challenging
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825. Cosmological Parameters

problem to the statistical analysis. We attempt a statistical fit to the whole Universe with 6
to 12 parameters, but we might need to include hundreds of nuisance parameters, some of themhighly correlated with the cosmological parameters of interest (for example time-dependent galaxy
biasing could mimic the growth of mass fluctuations). Fortunately, there is some astrophysicalprior knowledge on these effects, and a small number of physically-motivated free parameters would
ideally be preferred in the cosmological parameter analysis.
25.3.1 Measures of the Hubble constant
In 1929, Edwin Hubble discovered the law of expansion of the Universe by measuring distances to nearby galaxies. The slope of the relation between the distance and recession velocity is defined to be the present-epoch Hubble constant, H0 . Astronomers argued for decades
about the systematic uncertainties in various methods and derived values over the wide range< H < 100 km s−1 Mpc−1 .
40 km s−1 Mpc−1 ∼0 ∼One of the most reliable results on the Hubble constant came from the Hubble Space Telescope(HST) Key Project [20]. This study used the empirical period–luminosity relation for Cepheidvariable stars, and calibrated a number of secondary distance indicators: Type Ia Supernovae (SNeIa); the Tully–Fisher relation; surface-brightness fluctuations; and Type II Supernovae. Various
systematics are under investigation, for example JWST is helping understanding of the effect offield crowding. This approach has been further extended, e.g. using HST observations of Cepheids
in the hosts of 42 SNe Ia and exploiting Gaia EDR3 parallaxes, the SH0ES team derived H0 =(73.0 ± 1.0) km s−1 Mpc−1 [21].
Three other methods have been used recently. One is a calibration of the tip of the red-giant
branch applied to Type Ia supernovae, the Carnegie–Chicago Hubble Programme (CCHP), finding
H0 = (69.8 ± 0.6 (stat.) ± 1.6 (sys.)) km s−1 Mpc−1 [22]. The second uses the method of time
+4.1
delay in gravitationally-lensed quasars; Birrer et al. [23] find H0 = 67.4−3.2
km s−1 Mpc−1 froma sample of 40 lenses. A third method that came to fruition recently is based on gravitational
waves; the ‘bright standard siren’ method applied to the binary neutron star GW170817 yields−1 Mpc−1 [24]. Adding two ‘dark standard siren’ systems shifts this to H =H0 = 70+120−8 km s−1 Mpc−1 [25], still dominated by the single bright siren system. The 47 sources in the
73+11
km
s
−8
third Gravitational-Wave Transient Catalog (GWTC-3), again combined with the bright siren,−1 Mpc−1 [26]. When many more gravitational-wave events
yield the similar result H0 = 68+12−8 km shave been acquired, the future uncertainties on H0 from standard sirens will get smaller.
The determination of H0 by the Planck Collaboration [2] gives a lower value than most of theabove methods, H0 = (67.4 ± 0.5) km s−1 Mpc−1 . As they discuss, there is a strong degeneracyof H0 with other parameters, particularly Ωm and the neutrino mass. It is worth noting thatusing the ‘inverse distance ladder’ method gives a result H0 = (67.8 ± 1.3) km s−1 Mpc−1 [27],close to the Planck result. The inverse distance ladder relies on absolute-distance measurementsfrom baryon acoustic oscillations (BAOs) to calibrate the intrinsic magnitude of the SNe Ia (rather
than nearby Cepheids and parallax). This measurement was derived from 207 spectroscopicallyconfirmed Type Ia supernovae from the Dark Energy Survey (DES), an additional 122 low-redshift
SNe Ia, and measurements of BAOs. A combination of DES Year 3 (Y3) clustering and weaklensing with BAO and BBN (assuming ΛCDM) gives H0 = (67.6 ± 0.9) km s−1 Mpc−1 [28]. The
completed Extended Baryon Oscillation Spectroscopic Survey (eBOSS) [29] inverse distance ladderresult, within an assumed extended cosmological model, is H0 = (68.2 ± 0.8) km s−1 Mpc−1 , also
close to the Planck value.
The tension between the H0 values from Planck and the traditional cosmic distance-laddermethods, with the SH0ES result deviating from Planck by about 5σ, is under intense investigation
for potential systematic effects. There is possibly a trend for higher H0 derived from the nearby
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25. Cosmological Parameters

9

Universe and a lower H0 from the early Universe, which has led some researchers to propose a timevariation of the dark energy component or other exotic scenarios. Ongoing studies are addressing
the question of whether the Hubble tension is due to systematics in at least one of the probes, ora signature of new physics. Figure 25.1 shows a selection of recent H0 values, summarizing thecurrent status of the Hubble constant tension. See Ref. [30] for a review.
Planck 2018 – CMB

DES Year 3 – 3x2pt+BAO+BBN

SH0ES 2022 – Cepheids

CCHP – TRGBTDCOSMO + SLACS – Time delays

60

65

70
H0 [km s−1 Mpc−1 ]7580Figure 25.1: A selection of recent H0 measurements from the various projects as described in
the text. The standard-siren determinations are omitted as they are too wide for the plot. Figure
courtesy of Pablo Lemos.

25.3.2 Supernovae as cosmological probes
Empirically, the peak luminosity of SNe Ia can be used as an efficient distance indicator(e.g., Ref. [31]), thus allowing cosmology to be constrained via the distance–redshift relation. The
favorite theoretical explanation for SNe Ia is the thermonuclear disruption of carbon–oxygen whitedwarfs. Although not perfect ‘standard candles,’ it has been demonstrated that by correcting fora relation between the light-curve shape, color, and luminosity at maximum brightness, the dispersion of the measured luminosities can be greatly reduced. There are several possible systematiceffects that may affect the accuracy of the use of SNe Ia as distance indicators, e.g., evolution with
redshift and interstellar extinction in the host galaxy and in the Milky Way.
In the late 1990s two major studies, the Supernova Cosmology Project and the High-z Supernova
Search Team, found evidence for an accelerating Universe [32], interpreted as due to a cosmologicalconstant or a dark energy component. When combined with the CMB data (which indicate nearflatness, i.e., Ωm + ΩΛ ' 1), the best-fit values were Ωm ' 0.3 and ΩΛ ' 0.7. Most results inthe literature are consistent with the w = −1 cosmological constant case. The leading sample
currently is the Pantheon+ compilation [33]. This set of 1550 spectroscopically-confirmed SNe Ia
gives Ωm = 0.334 ± 0.018 (stat+sym) for an assumed flat ΛCDM model, while in combination with
+0.2
the CMB, for a flat wCDM model these data give w = −0.98−0.3[34]. Future experiments (e.g.
DES year 5 SNIa) will refine constraints on the cosmic equation of state w(z).
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1025. Cosmological Parameters25.3.3 Cosmic microwave background
The physics of the CMB is described in detail in the CMB chapter in this volume. Before
recombination, the baryons and photons are tightly coupled, and the perturbations oscillate in
the potential wells generated primarily by the dark matter perturbations. After decoupling, thebaryons are free to collapse into those potential wells. The CMB carries a record of conditions at
the time of last scattering, often called primary anisotropies. In addition, it is affected by various
processes as it propagates towards us, including the effect of a time-varying gravitational potential(the integrated Sachs-Wolfe effect), gravitational lensing, and scattering from ionized gas at low
redshift.
The primary anisotropies, the integrated Sachs–Wolfe effect, and the scattering from a homogeneous distribution of ionized gas, can all be calculated using linear perturbation theory. Available
codes include CAMB and CLASS [9], the former widely used embedded within the analysis package
CosmoMC [35] and in higher-level analysis packages such as CosmoSIS [36], CosmoLike [37], andCobaya [38]. Gravitational lensing is also calculated in these codes. Secondary effects, such as
inhomogeneities in the reionization process, and scattering from gravitationally-collapsed gas (theSunyaev–Zeldovich or SZ effect), require more complicated, and more uncertain, calculations.
The upshot is that the detailed pattern of anisotropies depends on all of the cosmologicalparameters. In a typical cosmology, the anisotropy power spectrum [usually plotted as `(`+1)C` /2π]
features a flat plateau at large angular scales (small `), followed by a series of oscillatory featuresat higher angular scales, the first and most prominent being at around one degree (` ' 200). Thesefeatures, known as acoustic peaks, represent the oscillations of the photon–baryon fluid aroundthe time of decoupling. Some features can be closely related to specific parameters—for instance,
the location in multipole space of the set of peaks probes the spatial geometry, while the relative
heights of the peaks probe the baryon density—but many other parameters combine to determine
the overall shape.
The 2018 data release from the Planck satellite [1] gives the most powerful results to date on
the spectrum of CMB temperature anisotropies, with a precision determination of the temperature
power spectrum to beyond ` = 2000. The Atacama Cosmology Telescope (ACT) and South Pole
Telescope (SPT) experiments extend these results to higher angular resolution, though without fullsky coverage. Planck and the polarization-sensitive versions of ACT and SPT give the current stateof the art in measuring the spectrum of E-polarization anisotropies and the correlation spectrum
between temperature and polarization. These are consistent with models based on the parameters
we have described, and provide accurate determinations of many of those parameters [2]. Primordial
B-mode polarization has not been detected (although the gravitational lensing effect on B modes
has been measured).
The data provide an exquisite measurement of the location of the set of acoustic peaks, determining the angular-diameter distance of the last-scattering surface. In combination with otherdata this strongly constrains the spatial geometry, in a manner consistent with spatial flatness andexcluding significantly-curved Universes. CMB data give a precision measurement of the age of the
Universe. The CMB also gives a baryon density consistent with, and at higher precision than, thatcoming from BBN. It affirms the need for both dark matter and dark energy. It shows no evidence
for dynamics of the dark energy, being consistent with a pure cosmological constant (w = −1).
The density perturbations are consistent with a power-law primordial spectrum, and there is no
indication yet of tensor perturbations. The current best-fit for the reionization optical depth from
CMB data, τ = 0.054, is in line with models of how early structure formation induces reionization.
Planck also made the first all-sky map of the CMB lensing field, which probes the entire matter
distribution in the Universe and adds some additional constraining power to the CMB-only datasets.
The recent ACT results [39] demonstrate very well the reconstruction of the mass power spectrum
31st May, 2024

1125. Cosmological Parametersat intermediate redshifts from CMB gravitational lensing. These measurements are consistent with
Planck, supporting ΛCDM (see more in the CMB chapter).
25.3.4 Galaxy clustering
The power spectrum of density perturbations is affected by the nature of the dark matter.
Within the ΛCDM model, the power spectrum shape depends primarily on the primordial power
spectrum and on the combination Ωm h, which determines the horizon scale at matter–radiation
equality, with a subdominant dependence on the baryon density. The matter distribution is mosteasily probed by observing the galaxy distribution, but this must be done with care since the galaxies
do not perfectly trace the dark matter distribution. Rather, they are a ‘biased’ tracer of the dark
matter [40]. The need to allow for such bias is emphasized by the observation that different types ofgalaxies show bias with respect to each other. In particular, scale-dependent and stochastic biasing
may introduce a systematic effect on the determination of cosmological parameters from redshift
surveys [41]. Prior knowledge from simulations of galaxy formation or from gravitational lensing
data could help to quantify biasing. Furthermore, the observed 3D galaxy distribution is in redshift
space, i.e., the observed redshift is the sum of the Hubble expansion and the line-of-sight peculiar
velocity, leading to linear and non-linear dynamical effects that also depend on the cosmologicalparameters. On the largest length scales, the galaxies are expected to trace the location of thedark matter, except for a constant multiplier b to the power spectrum, known as the linear bias
parameter. On scales smaller than 20 Mpc or so, the clustering pattern is ‘squashed’ in the radial
direction due to coherent infall, which depends approximately on the parameter β ≡ Ω0.6
m /b (onthese shorter scales, more complicated forms of biasing are not excluded by the data). On scales ofa few Mpc, there is an effect of elongation along the line of sight (colloquially known as the ‘finger
of God’ effect) that depends on the galaxy velocity dispersion.
25.3.4.1 Baryon acoustic oscillations
The power spectra of the 2-degree Field (2dF) Galaxy Redshift Survey and the Sloan Digital
Sky Survey (SDSS) are well fit by a ΛCDM model and both surveys showed first evidence for baryon
acoustic oscillations (BAOs) [42, 43]. When eBOSS is combined with Planck, Pantheon Type Ia
Supernovae and other probes the result is wp = −1.020 ± 0.032 at the pivot redshift zp = 0.29 [29].
Similar results for w were obtained e.g. by the WiggleZ survey [44]. Preliminary results from the
DESI galaxy redshift survey [45] confirm the detection of the BAOs, in agreement with previous
surveys.
25.3.4.2 Redshift distortion
There is continuing interest in Kaiser’s ‘redshift distortion’ effect [46]. This distortion depends
on cosmological parameters via the perturbation growth rate in linear theory f (z) = d ln δ/d ln a 'Ωγm (z), where γ ' 0.55 for the ΛCDM model and may be different for modified gravity models. By
measuring f (z) it is feasible to constrain γ and rule out certain modified gravity models [47,48]. We
note the degeneracy of the redshift-distortion pattern and the geometric distortion (the so-called
Alcock–Paczynski effect [49]), e.g., as illustrated by the WiggleZ survey [50] and eBOSS [29].
25.3.4.3 Limits on neutrino mass from galaxy surveys and other probes
Large-scale structure data place constraints on Ων due to the neutrino free-streaming effect [51].
Presently there is no clear detection, and upper limits on neutrino mass are commonly estimatedby comparing the observed galaxy power spectrum with a four-component model of baryons, cold
dark matter, a cosmological constant, and massive neutrinos. Such analyses also assume that theprimordial power spectrum is adiabatic, scale-invariant, and Gaussian. Potential systematic effects
include biasing of the galaxy distribution and non-linearities of the power spectrum. An upper limit
can also be derived from CMB anisotropies alone, while combination with additional cosmological31st May, 2024

1225. Cosmological Parameters

datasets can improve the results.
The most recent results on neutrino mass upper limits and other neutrino properties are summarized in the Neutrinos in Cosmology chapter in this volume. The latest cosmological data constrain
the sum of neutrino masses to be below 0.13 eV [28], consistent with Refs. [2, 29, 39]. Since thelower limit on this sum from oscillation experiments (assuming normal hierarchy) is 0.06 eV it isexpected that future cosmological surveys will soon detect effects from neutrino mass. Also, current
cosmological datasets are in good agreement with the standard value for the effective number ofneutrino species Neff = 3.044.
25.3.5 Clustering in the inter-galactic medium
It is commonly assumed, based on hydrodynamic simulations, that the neutral hydrogen in theinter-galactic medium (IGM) can be related to the underlying mass distribution. It is then possible
to estimate the matter power spectrum from the absorption observed in quasar spectra, the so-called
Lyman-α forest. The usual procedure is to measure the power spectrum of the transmitted flux,
and then to infer the mass power spectrum. Photo-ionization heating by the ultraviolet backgroundradiation and adiabatic cooling by the expansion of the Universe combine to give a simple powerlaw relation between the gas temperature and the baryon density. It also follows that there is a
power-law relation between the optical depth τ and ρb . Therefore, the observed flux F = exp(−τ )is strongly correlated with ρb , which itself traces the mass density. The matter and flux power
spectra can be related by a biasing function that is calibrated from simulations. There are twovariants of Lyman-alpha analyses: 1-dimensional power spectra from individual lines-of-sight thatprobe small (∼Mpc) scales; and 3-dimensional Lyman-alpha BAO analyses that measure large-scalecorrelations (over ∼100 Mpc scales) using neighboring quasar lines-of-sight.
The latest BOSS and eBOSS datasets have provided measurements of the BAO scale both inthe Lyman-α absorption and in its cross-correlation with quasars at an effective redshift z = 2.3[29, 52]. The results agree within 1.5-σ with the Planck CMB flat-Λ estimation for BAO scale.
Such measurements will improve with the Dark Energy Spectroscopic Instrument (DESI) and otherfuture spectroscopic surveys. The Lyman-α flux power spectrum has also been used to constrain
the nature of dark matter, for example limiting the amount of warm dark matter [53].

25.3.6 Weak gravitational lensing
Images of background galaxies are distorted by the gravitational effect of mass variations along
the line of sight. Deep gravitational potential wells, such as galaxy clusters, generate ‘strong lensing’leading to arcs, arclets, and multiple images, while more moderate perturbations give rise to ‘weaklensing.’ Weak lensing is now widely used to measure the mass power spectrum in selected regionsof the sky (see Ref. [54] for reviews). Since the signal is weak, the image of deformed galaxy shapes(the ‘shear map’) must be analyzed statistically to measure the power spectrum, higher moments,
and cosmological parameters. There are various systematic effects in the interpretation of weaklensing, e.g., due to atmospheric distortions during observations, the redshift distribution of the
background galaxies (usually depending on the accuracy of photometric redshifts), the intrinsic
correlation of galaxy shapes, and non-linear modeling uncertainties.
Weak-lensing measurements from the Kilo-Degree Survey (KiDS) [55], the Subaru Hyper-Suprime-Cam (HSC) [56], and DES [57] have constrained the clumpiness parameter S8 ≡ σ8 (Ωm /0.3)0.5 .
Each of these surveys has yielded S8 lower than that derived from Planck. Figure 25.2 showsthe Ωm –S8 constraints derived from a joint KiDS/DES analysis [58] versus the CMB constraint
from Planck, where the tension is significantly reduced, to 1.7σ, as compared to previous analyses.
Results from weak lensing from DES, combined with other probes, are described in the next section.
31st May, 2024

25. Cosmological Parameters

13

0.90

anck
ES Y3 + KiDS-1000
DS-1000
ES Y3

S8

0.85

P lanck
DES Y3 + KiDS-1000
KiDS-1000
DES Y3

0.80

0.75

0.70
0.4

0.5

0.2

0.3

0.4

0.5

m

Figure25.2:andMarginalized
posteriorcontours(innerconfidence
95% confidence
mological parameters
f8 (left)(8 (right) withthe matter
density⌦m in68%
flat-⇤CDM.
Welevel,adoptoutera Hybrid
from DES Y3 (green)
andKiDS-1000(yellow)andconductajoint-survey
analysis
of
DES
Y3
+
KiDS-1000
level) in the Ωm –S8 plane. The plot shows KiDS weak lensing alone, DES Y3 weak lensing alone,mpared to a re-analysisof thePlanckCollaboration(2020)CMB.
CMBFigureobservations(blue)
using a common
set ofthe DES andDES and
KiDS
combined,and Planck
courtesyof Alexandra
Amon andalised posterior KiDS
contours
show
the
68%
(inner)and95%(outer)credibleintervals.
collaborations.

marginal, maximum
DES parameters in green and the KiDS parameters in yele 4, it is worth
not-Other probes
low. We can compare the constraints from the individual and25.3.7
MAP and mean isOther
ex- probesjoint-survey
analyses. Analysing the surveys independently
that have been used to constrain cosmological parameters, but that are not��������2 mocks
in
we
find:
presently competitive in terms of accuracy, are the integrated Sachs-Wolfe effect [59, 60], the numignificant skew
the or composition ofDES+0.43[61], andKiDS+0.54velocities that probe theberindensity
galaxyclusters
galaxypeculiar= 0.32= 0.65(14)IA0.88 .
potential projection
mass fluctuations in the localIA
Universe [62]. 0.37
DESnalising over a neuInprobes
a joint-survey
25.4 Bringing
togetheranalysis, we find IA reduces to accomabout the truth (seea reductionin (8 relative
to the
DES-only
preferred
Although it modate
contains two
ingredients—darkmatter anddarkenergy—whichhave not yet been
follows we quote
theKiDS
value,experiments,
and the sameeﬀect,model
but inreverse,for IAaccepted:verified by laboratorythe ΛCDM
is almost
universallyby cosmologists
the reader to as
Table
4 description of the present data. The approximate values of somethe bestof the key parameters
Joint h ' 0.70.+0.52+0.29
maximum marginalare Ωb ' 0.05, Ωc 'DES,
0.25,JointΩΛ =' 0.70,a HubbleKiDS,
constant
The .spatial
0.02and= 1.04 0.54(15)geometry is
0.58
IA
IA

very close to flat (and usually assumed to be precisely flat), and the initial perturbations Gaussian,This scale-invariant.
increases the oﬀset between the survey-preferred IA
adiabatic, and nearly
amplitude,
but theconstraints
formally
ss density
The most powerful data sourceis the
CMB, whichremain
on its own
supportsconsistent
all these main tenets.
withaHellinger
oﬀsetof
1.47f.
Astheintrinsicalignment
Values for some parameters, as given in Ref. [2], are reproduced in Table 25.1. These particulareutrino mass results
density
is known
toconstraints
depend on
factors(see the bydiscuspresume signal
a flat Universe.
Thearemany
somewhatstrengthened
adding additionalon (2020) we inveswith ⌃< a = 0.06eV,
by oscillation expers hierarchy (Capozzi
o be fairly insensitive

sion in Appendix C.2), we do not expect identical intrinsic
31sttwo
May, surveys,
2024alignment signals in thehence our use of independent IA parameters. That said, as DES and KiDS havebroadly similar redshift distributions and depths, with the to-25. Cosmological Parameters

14

datasets, BAO being shown in the Table as an example, though most of the constraining power
resides in the CMB data. Similar constraints at lower precision were previously obtained by the
WMAP collaboration, and more recently by SPT.
Table 25.1: Parameter constraints reproduced from Ref. [2] (Table 2,column 5), with some additional rounding. Both columns assume theΛCDM cosmology with a power-law initial spectrum, no tensors, spatial
flatness, a cosmological constant as dark energy, and the sum of neutrino
masses fixed to 0.06 eV. Above the line are the six parameter combinations
actually fit to the data (θMC is a measure of the sound horizon at lastscattering); those below the line are derived from these. The first column
uses Planck primary CMB data plus the Planck measurement of CMB
lensing. This column gives our present recommended values. The second
column adds in data from a compilation of BAO measurements describedin Ref. [2]. The perturbation amplitude ∆2R (denoted As in the original
paper) is specified at the scale 0.05 Mpc−1 . Uncertainties are shown at
68% confidence.
Planck TT,TE,EE+lowE+lensing

+BAO

Ωb h2
Ωc h2
100 θMCnsτln(1010 ∆2R )0.02237 ± 0.00015
0.1200 ± 0.0012
1.0409 ± 0.0003
0.965 ± 0.004
0.054 ± 0.007
3.044 ± 0.014

0.02242 ± 0.00014
0.1193 ± 0.0009
1.0410 ± 0.0003
0.966 ± 0.004
0.056 ± 0.007
3.047 ± 0.014

h
σ8
Ωm
ΩΛ

0.674 ± 0.005
0.811 ± 0.006
0.315 ± 0.007
0.685 ± 0.007

0.677 ± 0.004
0.810 ± 0.006
0.311 ± 0.006
0.689 ± 0.006

If the assumption of spatial flatness is lifted, it turns out that the primary CMB on its own constrains the spatial curvature fairly weakly, due to a parameter degeneracy in the angular-diameter
distance. However, inclusion of other data readily removes this degeneracy. Simply adding the
Planck lensing measurement, and with the assumption that the dark energy is a cosmologicalP
constant, yields a 68% confidence constraint on Ωtot ≡Ωi + ΩΛ = 1.011 ± 0.006 and further
adding BAO makes it 0.9993 ± 0.0019 [2]. Results of this type are normally taken as justifying the
restriction to flat cosmologies.
One derived parameter that is very robust is the age of the Universe, since there is a usefulcoincidence that for a flat Universe the position of the first peak is strongly correlated with theage. The CMB data give 13.797 ± 0.023 Gyr (assuming flatness). This is in good agreement withthe ages of the oldest globular clusters and with radioactive dating.
The baryon density Ωb is now measured with high accuracy from CMB data alone, and isconsistent with and more precise than the determination from BBN. The value quoted in the
Big-Bang Nucleosynthesis chapter in this volume is Ωb h2 = 0.0220 ± 0.0004.
While ΩΛ is measured to be non-zero with very high confidence, there is no evidence of evolutionof the dark energy density. As described in the Dark Energy chapter in this volume a combination
31st May, 2024

1525. Cosmological Parametersof CMB, SN, RSD, and BAO measurements, assuming a flat Universe, yield w = −1.03 ± 0.03 [28],consistent with the cosmological constant case w = −1. Allowing more complicated forms of darkenergy weakens the limits.
The data provide strong support for the main predictions of the simplest inflation models:spatial flatness and adiabatic, Gaussian, nearly scale-invariant density perturbations. But it isdisappointing that there is no sign of primordial gravitational waves; combining Planck and WMAP
with BICEP2/Keck Array BK18 data (plus BAO data to help constrain ns ) gives a 95% confidence
upper limit of r < 0.036 at the scale 0.05 Mpc−1 [63]. The density perturbation spectral index isclearly required to be less than one by current data, though the strength of that conclusion canweaken if additional parameters are included in the model fits.
Tests have been made for various types of non-Gaussianity, a particular example being a parameter fNL that measures a quadratic contribution to the perturbations. Various non-Gaussian shapes
are possible (see Ref. [64] for details), and current constraints on the popular ‘local’, ‘equilateral’,local = −1 ± 5,
and ‘orthogonal’ types (combining CMB temperature and polarization data) are fNL
equil
ortho
fNL = −26 ± 47, and fNL = −38 ± 24, respectively (these may appear weak, but prominent
non-Gaussianity requires the product fNL ∆R to be large, and ∆R is of order 10−5 ). Clearly noneof these give any indication of primordial non-Gaussianity.
While the above results come from the CMB alone, other probes are becoming competitive(especially when considering more complex cosmological models), and so the combination of datafrom different sources is of growing importance. We note that it has become fashionable to combine
probes at the level of power-spectrum data vectors, taking into account nuisance parameters ineach type of measurement. Discussions on ‘tension’ in resulting cosmological parameters dependon the statistical approaches used. Commonly the cosmology community works within the Bayesianframework, and assesses agreement amongst data sets with respect to a model via Bayesian evidence,essentially the denominator in Bayes’ theorem.
As an example of results, combining DES Y3 (position–position clustering, galaxy–galaxy
lensing, and weak lensing shear) with Planck, BAO, and RSD measurements from eBOSS, andType Ia supernovae from the Pantheon dataset compilation, has shown the datasets to be mutuallycompatible and yields very tight constraints on cosmological parameters: S8 ≡ σ8 (Ωm /0.3)0.5 =0.812 ± 0.008 and Ωm = 0.306+0.004−0.005 in ΛCDM, and w = −1.03 ± 0.03 in wCDM [28] matching
the constraint in Ref. [29]. The combined measurement of the Hubble constant within ΛCDM
−1 Mpc−1 , still leaving substantial tension with the SH0ES measurement
gives H0 = 68.0+0.4−0.3 km s
described earlier. Future analyses and the next generation of surveys will test for deviations from
ΛCDM, for example epoch-dependent w(z) and modifications to General Relativity.

25.5 Outlook for the future
The concordance model is now well-established, and there seems little room left for any dramatic
revision of this paradigm. A measure of the strength of that statement is how difficult it has provento formulate convincing alternatives.
Should there indeed be no major revision of the current paradigm, we can expect future developments to take one of two directions. Either the existing parameter set will continue to prove
sufficient to explain the data, with the parameters subject to ever-tightening constraints, or it
will become necessary to deploy new parameters. The latter outcome would be very much themore interesting, offering a route towards understanding new physical processes relevant to thecosmological evolution. There are many possibilities on offer for striking discoveries, for example:• the cosmological effects of a neutrino mass may be unambiguously detected, shedding lighton fundamental neutrino properties;
31st May, 2024

16

25. Cosmological Parameters• detection of primordial non-Gaussianities would indicate that non-linear processes influence
the perturbation generation mechanism;• detection of variation in the dark energy density (i.e., w 6= −1) would provide much-needed
experimental input into its nature.

These supply more than enough motivation for continued efforts to test the cosmological model andimprove its accuracy. Over the coming years, there are a wide range of new observations that willbring further precision to cosmological studies. Indeed, there are far too many to mention them all
here, and so we highlight just a few areas.
The CMB observations will improve in several directions. A current frontier is the study ofpolarization, for which power spectrum measurements have now been made by several experiments.
Detection of primordial B-mode anisotropies is the next major goal and a variety of projects aretargeting this, though theory gives little guidance as to the likely signal level. Future CMB projects
that are approved include the Simons Observatory and LiteBIRD.
An impressive array of cosmology surveys are already operational, under construction, or proposed, including the ground-based Hyper Suprime Camera (HSC) and Rubin-LSST imaging surveys, spectroscopic surveys such as DESI, and space missions Euclid (succesfully launched in July2023) and the Roman Space Telescope.
An exciting area for the future is radio surveys of the redshifted 21-cm line of hydrogen. Becauseof the intrinsic narrowness of this line, by tuning the bandpass the emission from narrow redshift
slices of the Universe will be measured to extremely high redshift, probing the details of the reionization process at redshifts up to perhaps 20, as well as measuring large-scale features such as the
BAOs. LOFAR and CHIME are the first instruments able to do this and have begun operations. In
the longer term, the Square Kilometre Array (SKA) will take these studies to sophisticated levels.
The development of the first precision cosmological model is a major achievement. However, it isimportant not to lose sight of the motivation for developing such a model, which is to understand theunderlying physical processes at work governing the Universe’s evolution. From that perspective,
progress has been much less dramatic. For instance, there are many proposals for the nature of
the dark matter, but no consensus as to which is correct. The nature of the dark energy remains amystery. Even the baryon density, now measured to an accuracy of a percent, lacks an underlying
theory able to predict it within orders of magnitude. Precision cosmology may have arrived, but
at present many key questions remain, to motivate and challenge the cosmology community.
References[1] Planck Collab. 2018 Results I, Astron. Astrophys. 641, A1 (2020), [arXiv:1807.06205].
[2] Planck Collab. 2018 Results VI, Astron. Astrophys. 641, A6 (2020), [arXiv:1807.06209].
[3] C. L. Bennett et al. (WMAP), Astrophys. J. Suppl. 208, 20 (2013), [arXiv:1212.5225].
[4] G. Hinshaw et al. (WMAP), Astrophys. J. Suppl. 208, 19 (2013), [arXiv:1212.5226].
[5] S. Fukuda et al. (Super-Kamiokande), Phys. Rev. Lett. 85, 3999 (2000), [hep-ex/0009001];
Q. R. Ahmad et al. (SNO), Phys. Rev. Lett. 87, 071301 (2001), [arXiv:nucl-ex/0106015].
[6] E.W. Kolb and M.S. Turner, The Early Universe, Addison–Wesley (Redwood City, 1990).
[7] D.H. Lyth and A.R. Liddle, The Primordial Density Perturbation, Cambridge University Press(2009).
[8] A. R. Liddle and D. H. Lyth, Phys. Lett. B291, 391 (1992), [arXiv:astro-ph/9208007].
[9] A. Lewis, A. Challinor and A. Lasenby, Astrophys. J. 538, 473 (2000), [arXiv:astroph/9911177]; D. Blas, J. Lesgourgues and T. Tram, JCAP 1107, 034 (2011), [arXiv:1104.2933].
[10] D. J. Fixsen, Astrophys. J. 707, 916 (2009), [arXiv:0911.1955].
31st May, 2024

17

25. Cosmological Parameters[11] M. Hobson et al. (eds). Bayesian Methods in Cosmology, Cambridge University Press (2009).
[12] A. Kosowsky and M. S. Turner, Phys. Rev. D52, R1739 (1995), [arXiv:astro-ph/9504071].
[13] K. A. Malik and D. Wands, Phys. Rept. 475, 1 (2009), [arXiv:0809.4944].
[14] Planck Collab. 2013 Results XXV, Astron. Astrophys. 571, A25 (2014), [arXiv:1303.5085].
[15] D. H. Lyth and D. Wands, Phys. Lett. B524, 5 (2002), [hep-ph/0110002]; K. Enqvist and
M. S. Sloth, Nucl. Phys. B626, 395 (2002), [hep-ph/0109214]; T. Moroi and T. Takahashi,
Phys. Lett. B522, 215 (2001), [Erratum: Phys. Lett.B539,303(2002)], [hep-ph/0110096].
[16] J. J. Bennett et al., JCAP 04, 073 (2021), [arXiv:2012.02726].
[17] S. Riemer-Sorensen, D. Parkinson and T. M. Davis, Publ. Astron. Soc. Austral. 30, e029
(2013), [arXiv:1301.7102].
[18] J. K. Webb et al., Phys. Rev. Lett. 107, 191101 (2011), [arXiv:1008.3907]; J. A. King et al.,Mon. Not. Roy. Astron. Soc. 422, 3370 (2012), [arXiv:1202.4758]; P. Molaro et al., Astron. &Astrophys. 555, 68 (2013).
[19] Planck Collab. 2015 Results XVIII, Astron. Astrophys. 594, A18 (2016), [arXiv:1502.01593].
[20] W. L. Freedman et al. (HST), Astrophys. J. 553, 47 (2001), [arXiv:astro-ph/0012376].
[21] A. G. Riess et al., Astrophys. J. Lett. 934, 1, L7 (2022), [arXiv:2112.04510].
[22] W. L. Freedman, Astrophys. J. 919, 1, 16 (2021), [arXiv:2106.15656].
[23] S. Birrer et al., Astron. Astrophys. 643, A165 (2020), [arXiv:2007.02941].
[24] B. P. Abbott et al. (LIGO Scientific, Virgo, 1M2H, Dark Energy Camera GW-E, DES,
DLT40, Las Cumbres Observatory, VINROUGE, MASTER), Nature 551, 7678, 85 (2017),[arXiv:1710.05835].
[25] A. Palmese et al., Astrophys. J. 943, 1, 56 (2023), [arXiv:2111.06445].
[26] R. Abbott et al. (LIGO Scientific, Virgo„ KAGRA, VIRGO), Astrophys. J. 949, 2, 76 (2023),[arXiv:2111.03604].
[27] E. Macaulay et al. (DES), Mon. Not. Roy. Astron. Soc. 486, 2, 2184 (2019), [arXiv:1811.02376].
[28] T. M. C. Abbott et al. (DES), Phys. Rev. D 105, 2, 023520 (2022), [arXiv:2105.13549].
[29] S. Alam et al. (eBOSS), Phys. Rev. D 103, 8, 083533 (2021), [arXiv:2007.08991].
[30] P. Shah, P. Lemos and O. Lahav, Astron. Astrophys. Rev. 29, 1, 9 (2021), [arXiv:2109.01161].
[31] B. Leibundgut, Ann. Rev. Astron. Astrophys. 39, 67 (2001).
[32] A. G. Riess et al. (Supernova Search Team), Astron. J. 116, 1009 (1998), [arXiv:astroph/9805201]; P. M. Garnavich et al. (Supernova Search Team), Astrophys. J. 509, 74 (1998),[arXiv:astro-ph/9806396]; S. Perlmutter et al. (Supernova Cosmology Project), Astrophys. J.
517, 565 (1999), [arXiv:astro-ph/9812133].
[33] D. Scolnic et al., Astrophys. J. 938, 2, 113 (2022), [arXiv:2112.03863].
[34] D. Brout et al., Astrophys. J. 938, 2, 110 (2022), [arXiv:2202.04077].
[35] A. Lewis and S. Bridle, Phys. Rev. D66, 103511 (2002), [arXiv:astro-ph/0205436].
[36] J. Zuntz et al., Astron. Comput. 12, 45 (2015), [arXiv:1409.3409].
[37] E. Krause and T. Eifler, Mon. Not. Roy. Astron. Soc. 470, 2, 2100 (2017), [arXiv:1601.05779].
[38] J. Torrado and A. Lewis, JCAP 05, 057 (2021), [arXiv:2005.05290].
[39] M. S. Madhavacheril et al. (ACT) (2023), [arXiv:2304.05203].
[40] N. Kaiser, Astrophys. J. 284, L9 (1984).
31st May, 2024

18

25. Cosmological Parameters[41] A. Dekel and O. Lahav, Astrophys. J. 520, 24 (1999), [arXiv:astro-ph/9806193].
[42] D. J. Eisenstein et al. (SDSS), Astrophys. J. 633, 560 (2005), [arXiv:astro-ph/0501171].
[43] S. Cole et al. (2dFGRS), Mon. Not. Roy. Astron. Soc. 362, 505 (2005), [arXiv:astroph/0501174].
[44] D. Parkinson et al., Phys. Rev. D86, 103518 (2012), [arXiv:1210.2130].
[45] J. Moon et al. (DESI) (2023), [arXiv:2304.08427].
[46] N. Kaiser, Mon. Not. Roy. Astron. Soc. 227, 1 (1987).
[47] L. Guzzo et al., Nature 451, 541 (2008), [arXiv:0802.1944].
[48] A. Nusser and M. Davis, Astrophys. J. 736, 93 (2011), [arXiv:1101.1650].
[49] C. Alcock and B. Paczynski, Nature 281, 358 (1979).
[50] C. Blake et al., Mon. Not. Roy. Astron. Soc. 425, 405 (2012), [arXiv:1204.3674].
[51] J. Lesgourgues and S. Pastor, Phys. Rept. 429, 307 (2006), [arXiv:astro-ph/0603494].
[52] H. du Mas des Bourboux et al., Astrophys. J. 901, 2, 153 (2020), [arXiv:2007.08995].
[53] M. Viel et al., Phys. Rev. D88, 043502 (2013), [arXiv:1306.2314].
[54] A. Refregier, Ann. Rev. Astron. Astrophys. 41, 645 (2003), [arXiv:astro-ph/0307212];
R. Massey et al., Nature 445, 286 (2007), [arXiv:astro-ph/0701594]; H. Hoekstra and B. Jain,
Ann. Rev. Nucl. Part. Sci. 58, 99 (2008), [arXiv:0805.0139].
[55] M. Asgari et al. (KiDS), Astron. Astrophys. 645, A104 (2021), [arXiv:2007.15633].
[56] C. Hikage et al. (HSC), Publ. Astron. Soc. Jap. 71, 2, 43 (2019), [arXiv:1809.09148].
[57] L. F. Secco et al. (DES) (2021), [arXiv:2105.13544].
[58] T. M. C. Abbott et al. (Kilo-Degree Survey, DES) (2023), [arXiv:2305.17173].
[59] R. G. Crittenden and N. Turok, Phys. Rev. Lett. 75, 2642 (1995), [arXiv:astro-ph/9505120].
[60] Planck Collab. 2015 Results XXI, Astron. Astrophys. 594, A21 (2016), [arXiv:1502.01595].
[61] Planck Collab. 2015 Results XXIV, Astron. Astrophys. 594, A24 (2016), [arXiv:1502.01597].
[62] A. Dekel, Ann. Rev. Astron. Astrophys. 32, 371 (1994), [arXiv:astro-ph/9401022].
[63] P. A. R. Ade et al. (BICEP/Keck), Phys. Rev. Lett. 127, 15, 151301 (2021), [arXiv:2110.00483].
[64] Planck Collab. 2018 Results IX, Astron. Astrophys. 641, A9 (2020), [arXiv:1905.05697].

31st May, 2024