Friedmann equations#Density parameter

{{main|Friedmann–Lemaître–Robertson–Walker metric}}

The Friedmann equations build on three assumptions:{{rp|22.1.3}}

1. the Friedmann–Lemaître–Robertson–Walker metric,
2. Einstein's equations for general relativity, and
3. a perfect fluid source.

The metric in turn starts with the simplifying assumption that the universe is spatially homogeneous and isotropic, that is, the cosmological principle; empirically, this is justified on scales larger than the order of 100 Mpc.

The metric can be written as:{{Cite book |last=Peacock |first=J. A. |url=https://www.cambridge.org/core/product/identifier/9780511804533/type/book |title=Cosmological Physics |date=1998-12-28 |publisher=Cambridge University Press |isbn=978-0-521-41072-4 |edition=1 |doi=10.1017/cbo9780511804533}}{{rp|65}}

$$c^2d\backslash tau^2\; =\; c^2dt^2\; -\; R^2(t)\; \backslash left(dr^2\; +\; S^2\_k(r)\; d\backslash psi^2\backslash right)$$

where

$$S\_\{-1\}(r)\; =\; \backslash sinh(r),\; S\_0\; =\; 1,\; S\_1\; =\; \backslash sin(r).$$

These three possibilities correspond to parameter {{mvar|k}} of (0) flat space, (+1) a sphere of constant positive curvature or (−1) a hyperbolic space with constant negative curvature.

Here the radial position has been decomposed into a time-dependent scale factor, $R(t)$, and a comoving coordinate, $r$.

Inserting this metric into Einstein's field equations relate the evolution of this scale factor to the pressure and energy of the matter in the universe. With the stress–energy tensor for a perfect fluid, results in the equations are described below.{{rp|p=73}}

{{General relativity sidebar |equations}}

There are two independent Friedmann equations for modelling a homogeneous, isotropic universe.

The first is:{{Cite journal |last=Navas, S. |year=2024 |title=Review of Particle Physics |journal=Physical Review D |volume=110 |issue=3 |pages=1–708 |doi=10.1103/PhysRevD.110.030001 |collaboration=Particle Data Group|hdl=20.500.11850/695340 |hdl-access=free }} 22.1.3 The Friedmann equations of motion

$$H^2\backslash equiv\; \{\backslash left(\backslash frac\{\backslash dot\{R\}\}\{R\}\backslash right)\}^2\; =\; \backslash frac\{8\; \backslash pi\; G\; \backslash rho\}\{3\}\; -\; \backslash frac\{k\}\{R^2\}\; +\; \backslash frac\{\backslash Lambda\}\{3\},$$ and second is:

$$\backslash frac\{\backslash ddot\{R\}\}\{R\}\; =\; \backslash frac\{\backslash Lambda\}\{3\}\; -\backslash frac\{4\; \backslash pi\; G\}\{3\}\backslash left(\backslash rho+3p\backslash right).$$

The term Friedmann equation sometimes is used only for the first equation.

In these equations,

{{mvar|R(t)}} is the cosmological scale factor, $G$ is the Newtonian constant of gravitation, {{math|Λ}} is the cosmological constant with dimension length⁻², {{mvar|ρ}} is the energy density and {{mvar|p}} is the isotropic pressure. {{mvar|k}} is constant throughout a particular solution, but may vary from one solution to another. The units set the speed of light in vacuum to one.

In previous equations, {{mvar|R}}, {{mvar|ρ}}, and {{mvar|p}} are functions of time. If the cosmological constant, {{math|Λ}}, is ignored, the term $-k/R^2$ in the first Friedmann equation can be interpreted as a Newtonian total energy, so the evolution of the universe pits gravitational potential energy, $8\backslash pi\; G\backslash rho/3$ against kinetic energy, $\backslash dot\{R\}/R$. The winner depends upon the {{mvar|k}} value in the total energy: if k is +1, gravity eventually causes the universe to contract. These conclusions will be altered if the {{math|Λ}} is not zero.

Using the first equation, the second equation can be re-expressed as:

$$\backslash dot\{\backslash rho\}\; =\; -3\; H\; \backslash left(\backslash rho\; +\; \backslash frac\{p\}\{c^2\}\backslash right),$$

which eliminates {{math|Λ}}. Alternatively the conservation of mass–energy:

$$T^\{\backslash alpha\backslash beta\}\{\}\_\{;\backslash beta\}=\; 0$$

leads to the same result.

The first Friedmann equation contains a discrete parameter {{math|1=k = +1, 0}} or {{math|−1}} depending on whether the shape of the universe is a closed 3-sphere, flat (Euclidean space) or an open 3-hyperboloid, respectively.{{Cite book |last=D'Inverno |first=Ray |title=Introducing Einstein's relativity |date=2008 |publisher=Clarendon Press |isbn=978-0-19-859686-8 |edition=Repr |location=Oxford}} If {{mvar|k}} is positive, then the universe is "closed": starting off on some paths through the universe return to the starting point. Such a universe is analogous to a sphere: finite but unbounded. If {{mvar|k}} is negative, then the universe is "open": infinite and no paths return. If {{math|1=k = 0}}, then the universe is Euclidean (flat) and infinite.{{rp|69}}

A dimensionless scale factor can be defined:

$$a(t)\; \backslash equiv\; \backslash frac\{R(t)\}\{R\_0\}$$

using the present day value

$$R\_0\; =\; R(\backslash text\{now\}).$$

The Friedmann equations can be written in terms of this dimensionless scale factor:

$$H^2(t)\; =\; \backslash left(\; \backslash frac\{\backslash dot\{a\}\}\{a\}\backslash right)^2\; =\; \backslash frac\{8\backslash pi\; G\}\{3\}\backslash left[\; \backslash rho(t)\; +\; \backslash frac\{\backslash rho\_c\; -\; \backslash rho\_0\}\{a^2(t)\}\backslash right]$$

where $\backslash dot\{a\}\; =\; da/dt$, $\backslash rho\_c\; =\; 3H^2\_0/8\backslash pi\; G$, and $\backslash rho\_0\; =\; \backslash rho(t=\backslash text\{now\})$.{{Cite book |last=Dodelson |first=Scott |title=Modern cosmology |date=2003 |publisher=Academic Press |isbn=978-0-12-219141-1 |location=San Diego, Calif}}{{rp|3}}

That value of the mass-energy density, $\backslash rho$ that gives $k=0$ when $\backslash Lambda=0$ is called the critical density:

$$\backslash rho\_c\; \backslash equiv\; \backslash frac\{3H^2\}\{8\backslash pi\; G\}.$$

If the universe has higher density, $\backslash rho\; \backslash ge\; \backslash rho\_c$, then it is called "spatially closed": in this simple approximation the universe would eventually contract. On the other hand, if has lower density, $\backslash rho\; \backslash le\; \backslash rho\_c$, then it is called "spatially open" and expands forever. Therefore the geometry of the universe is directly connected to its density.{{rp|p=73}}

== Density parameter ==

The density parameter {{mvar|Ω}} is defined as the ratio of the actual (or observed) density {{mvar|ρ}} to the critical density {{math|ρ_c}} of the Friedmann universe:{{rp|p=74}}

$$\backslash Omega\; :=\; \backslash frac\{\backslash rho\}\{\backslash rho\_\{c\}\}\; =\; \backslash frac\{8\; \backslash pi\; G\backslash rho\}\{3\; H^2\}.$$

Both the density $\backslash rho(t)$ and the Hubble parameter $H(t)$ depend upon time and thus the density parameter varies with time.{{rp|p=74}}

The critical density is equivalent to approximately five atoms (of monatomic hydrogen) per cubic metre, whereas the average density of ordinary matter in the Universe is believed to be 0.2–0.25 atoms per cubic metre.{{Cite book |last=Rees |first=Martin |title=Just six numbers: the deep forces that shape the universe |date=2001 |publisher=Basic Books |isbn=978-0-465-03673-8 |edition=Repr. |series=Astronomy/science |location=New York, NY}}{{cite web | publisher=NASA | title=Universe 101 | url=http://map.gsfc.nasa.gov/universe/uni_matter.html | access-date=September 9, 2015 | quote=The actual density of atoms is equivalent to roughly 1 proton per 4 cubic meters.}}

File:UniverseComposition.svg dominates the total energy (74%) while dark matter (22%) constitutes most of the mass. Of the remaining baryonic matter (4%), only one tenth is compact. In February 2015, the European-led research team behind the Planck cosmology probe released new data refining these values to 4.9% ordinary matter, 25.9% dark matter and 69.1% dark energy.]]

A much greater density comes from the unidentified dark matter, although both ordinary and dark matter contribute in favour of contraction of the universe. However, the largest part comes from so-called dark energy, which accounts for the cosmological constant term. Although the total density is equal to the critical density (exactly, up to measurement error), dark energy does not lead to contraction of the universe but rather may accelerate its expansion.

An expression for the critical density is found by assuming {{mvar|Λ}} to be zero (as it is for all basic Friedmann universes) and setting the normalised spatial curvature, {{mvar|k}}, equal to zero. When the substitutions are applied to the first of the Friedmann equations given the new $H\_0$ value we find:{{Cite journal |last1=Scolnic |first1=Daniel |last2=Riess |first2=Adam G. |last3=Murakami |first3=Yukei S. |last4=Peterson |first4=Erik R. |last5=Brout |first5=Dillon |last6=Acevedo |first6=Maria |last7=Carreres |first7=Bastien |last8=Jones |first8=David O. |last9=Said |first9=Khaled |last10=Howlett |first10=Cullan |last11=Anand |first11=Gagandeep S. |date=2025-01-15 |title=The Hubble Tension in Our Own Backyard: DESI and the Nearness of the Coma Cluster |journal=The Astrophysical Journal Letters |volume=979 |issue=1 |pages=L9 |doi=10.3847/2041-8213/ada0bd |doi-access=free |bibcode=2025ApJ...979L...9S |issn=2041-8205|arxiv=2409.14546 }}

$$\backslash begin\{align\}$$

\rho = \frac{3 H_0^2}{8 \pi G}

&\approx 1.10 \times 10^{-26} \mathrm{kg \, m^{-3}}

\\&\approx 1.88 \times 10^{-26} {\rm h}^2\, {\rm kg}\, {\rm m}^{-3}

\\&\approx 2.78 \times 10^{11} h^2 M_\odot\,{\rm Mpc}^{-3}

\end{align}

where:

• $H\_0\; =\; 76.5\; \backslash pm\; 2.2\; \backslash ,\; \backslash mathrm\{km\; \backslash ,\; s^\{-1\}\; \backslash ,\; Mpc^\{-1\}\}$

\approx 2.48 \times 10^{-18} \mathrm{s^{-1}}

• $h\; =\; \backslash frac\{H\_0\}\{100\; \backslash ,\; \backslash mathrm\{(km/s)/Mpc\}\}$
• $\backslash rho\_c\; =\; 8.5\; \backslash times\; 10^\{-27\}\; \backslash mathrm\{kg\; /\; m^3\}$

Given the value of dark energy to be $\backslash Omega\_\backslash Lambda\; =\; 0.647$

This term originally was used as a means to determine the spatial geometry of the universe, where {{math|ρ_c}} is the critical density for which the spatial geometry is flat (or Euclidean). Assuming a zero vacuum energy density, if {{mvar|Ω}} is larger than unity, the space sections of the universe are closed; the universe will eventually stop expanding, then collapse. If {{mvar|Ω}} is less than unity, they are open; and the universe expands forever. However, one can also subsume the spatial curvature and vacuum energy terms into a more general expression for {{mvar|Ω}} in which case this density parameter equals exactly unity. Then it is a matter of measuring the different components, usually designated by subscripts. According to the ΛCDM model, there are important components of {{mvar|Ω}} due to baryons, cold dark matter and dark energy. The spatial geometry of the universe has been measured by the WMAP spacecraft to be nearly flat. This means that the universe can be well approximated by a model where the spatial curvature parameter {{mvar|k}} is zero; however, this does not necessarily imply that the universe is infinite: it might merely be that the universe is much larger than the part we see.

The first Friedmann equation is often seen in terms of the present values of the density parameters, that is{{cite journal | last=Nemiroff | first=Robert J. | author-link=Robert J. Nemiroff | author2=Patla, Bijunath |arxiv = astro-ph/0703739| doi = 10.1119/1.2830536 | volume=76 | title=Adventures in Friedmann cosmology: A detailed expansion of the cosmological Friedmann equations | journal=American Journal of Physics | year=2008 | issue=3 | pages=265–276 | bibcode = 2008AmJPh..76..265N| s2cid=51782808 }}

$$\backslash frac\{H^2\}\{H\_0^2\}\; =\; \backslash Omega\_\{0,\backslash mathrm\; R\}\; a^\{-4\}\; +\; \backslash Omega\_\{0,\backslash mathrm\; M\}\; a^\{-3\}\; +\; \backslash Omega\_\{0,k\}\; a^\{-2\}\; +\; \backslash Omega\_\{0,\backslash Lambda\}.$$

Here {{math|Ω_0,R}} is the radiation density today (when {{math|1=a = 1}}), {{math|Ω_0,M}} is the matter (dark plus baryonic) density today, {{math|1=Ω_0,k = 1 − Ω₀}} is the "spatial curvature density" today, and {{math|Ω_0,Λ}} is the cosmological constant or vacuum density today.

{{Unreferenced section|date=September 2024}}

The Hubble parameter can change over time if other parts of the equation are time dependent (in particular the mass density, the vacuum energy, or the spatial curvature). Evaluating the Hubble parameter at the present time yields Hubble's constant which is the proportionality constant of Hubble's law. Applied to a fluid with a given equation of state, the Friedmann equations yield the time evolution and geometry of the universe as a function of the fluid density.

Relativisitic cosmology models based on the FLRW metric and obeying the Friedmann equations are called FRW models.{{rp|p=73}}

Direct observation of stars has shown their velocities to be dominated by radial recession, validating these assumptions for cosmological models.{{rp|p=65}}

These models are the basis of the standard model{{Cite book |last1=Bergström |first1=Lars |url=https://books.google.com/books?id=CQYu_sutWAoC&pg=PA61 |title=Cosmology and particle astrophysics |last2=Goobar |first2=Ariel |date=2008 |publisher=Praxis Publ |isbn=978-3-540-32924-4 |edition=2. ed., reprinted |series=Springer Praxis books in astronomy and planetary science |location=Chichester, UK |page=61}} of Big Bang cosmological including the current ΛCDM model.{{rp|loc=25.1.3}}

To apply the metric to cosmology and predict its time evolution via the scale factor $a(t)$ requires Einstein's field equations together with a way of calculating the density, $\backslash rho\; (t),$ such as a cosmological equation of state.

This process allows an approximate analytic solution Einstein's field equations $G\_\{\backslash mu\backslash nu\}\; +\; \backslash Lambda\; g\_\{\backslash mu\backslash nu\}\; =\; \backslash kappa\; T\_\{\backslash mu\backslash nu\}$ giving the Friedmann equations when the energy–momentum tensor is similarly assumed to be isotropic and homogeneous. The resulting equations are:{{Cite journal |last1=Rosu |first1=H. C. |last2=Ojeda-May |first2=P. |date=June 2006 |title=Supersymmetry of FRW Barotropic Cosmologies |journal=International Journal of Theoretical Physics |language=en |volume=45 |issue=6 |pages=1152–1157 |arxiv=gr-qc/0510004 |bibcode=2006IJTP...45.1152R |doi=10.1007/s10773-006-9123-2 |issn=0020-7748 |s2cid=119496918}}

$$\backslash begin\{align\}$$

{\left(\frac{\dot a}{a}\right)}^2 + \frac{kc^2}{a^2} - \frac{\Lambda c^2}{3} &= \frac{\kappa c^4}{3}\rho \\[4pt]

2\frac{\ddot a}{a} + {\left(\frac{\dot a}{a}\right)}^2 + \frac{kc^2}{a^2} - \Lambda c^2 &= -\kappa c^2 p .

\end{align}

Because the FLRW model assumes homogeneity, some popular accounts mistakenly assert that the Big Bang model cannot account for the observed lumpiness of the universe. In a strictly FLRW model, there are no clusters of galaxies or stars, since these are objects much denser than a typical part of the universe. Nonetheless, the FLRW model is used as a first approximation for the evolution of the real, lumpy universe because it is simple to calculate, and models that calculate the lumpiness in the universe are added onto the FLRW models as extensions. Most cosmologists agree that the observable universe is well approximated by an almost FLRW model, i.e., a model that follows the FLRW metric apart from primordial density fluctuations. {{As of|2003}}, the theoretical implications of the various extensions to the FLRW model appear to be well understood, and the goal is to make these consistent with observations from COBE and WMAP.

The pair of equations given above is equivalent to the following pair of equations

$$\backslash begin\{align\}$$

\dot{\rho} &= - 3 \frac{\dot a}{a} \left(\rho+\frac{p}{c^2}\right) \\[1ex]

\frac{\ddot{a}}{a} &= - \frac{\kappa c^4}{6} \left(\rho + \frac{3p}{c^2}\right) + \frac{\Lambda c^2}{3}

\end{align}

with $k$, the spatial curvature index, serving as a constant of integration for the first equation.

The first equation can be derived also from thermodynamical considerations and is equivalent to the first law of thermodynamics, assuming the expansion of the universe is an adiabatic process (which is implicitly assumed in the derivation of the Friedmann–Lemaître–Robertson–Walker metric).

The second equation states that both the energy density and the pressure cause the expansion rate of the universe $\{\backslash dot\; a\}$ to decrease, i.e., both cause a deceleration in the expansion of the universe. This is a consequence of gravitation, with pressure playing a similar role to that of energy (or mass) density, according to the principles of general relativity. The cosmological constant, on the other hand, causes an acceleration in the expansion of the universe.

The cosmological constant term can be omitted if we make the following replacements

$$\backslash begin\{align\}$$

\rho &\to \rho - \frac{\Lambda}{\kappa c^2}, &

p &\to p + \frac{\Lambda}{\kappa}.

\end{align}

Therefore, the cosmological constant can be interpreted as arising from a form of energy that has negative pressure, equal in magnitude to its (positive) energy density:

$$p\; =\; -\; \backslash rho\; c^2\; \backslash ,,$$

which is an equation of state of vacuum with dark energy.

An attempt to generalize this to

$$p\; =\; w\; \backslash rho\; c^2$$

would not have general invariance without further modification.

In fact, in order to get a term that causes an acceleration of the universe expansion, it is enough to have a scalar field that satisfies

Such a field is sometimes called quintessence.

This is due to McCrea and Milne,{{cite journal |last1=McCrea |first1=W. H. |last2=Milne |first2=E. A. |year=1934 |title=Newtonian universes and the curvature of space |journal=Quarterly Journal of Mathematics |volume=5 |pages=73–80 |doi=10.1093/qmath/os-5.1.73 |bibcode=1934QJMat...5...73M }} although sometimes incorrectly ascribed to Friedmann.{{according to whom|date=February 2025}} The Friedmann equations are equivalent to this pair of equations:

$$\backslash begin\{align\}$$

- a^3 \dot{\rho} = 3 a^2 \dot{a} \rho + \frac{3 a^2 p \dot{a}}{c^2} \, \\[1ex]

\frac{\dot{a}^2}{2} - \frac{\kappa c^4 a^3 \rho}{6a} = - \frac{k c^2}{2} \,.

\end{align}

The first equation says that the decrease in the mass contained in a fixed cube (whose side is momentarily a) is the amount that leaves through the sides due to the expansion of the universe plus the mass equivalent of the work done by pressure against the material being expelled. This is the conservation of mass–energy (first law of thermodynamics) contained within a part of the universe.

The second equation says that the kinetic energy (seen from the origin) of a particle of unit mass moving with the expansion plus its (negative) gravitational potential energy (relative to the mass contained in the sphere of matter closer to the origin) is equal to a constant related to the curvature of the universe. In other words, the energy (relative to the origin) of a co-moving particle in free-fall is conserved. General relativity merely adds a connection between the spatial curvature of the universe and the energy of such a particle: positive total energy implies negative curvature and negative total energy implies positive curvature.

The cosmological constant term is assumed to be treated as dark energy and thus merged into the density and pressure terms.

During the Planck epoch, one cannot neglect quantum effects. So they may cause a deviation from the Friedmann equations.

The Friedmann equations can be solved exactly in presence of a perfect fluid with equation of state

$$p\; =\; w\; \backslash rho\; c^2,$$

where {{mvar|p}} is the pressure, {{mvar|ρ}} is the mass density of the fluid in the comoving frame and {{mvar|w}} is some constant.

In spatially flat case ({{math|1=k = 0}}), the solution for the scale factor is

$$a(t)\; =\; a\_0\; \backslash ,\; t^\{\backslash frac\{2\}\{3(w+1)\}\}$$

where {{math|a₀}} is some integration constant to be fixed by the choice of initial conditions. This family of solutions labelled by {{mvar|w}} is extremely important for cosmology. For example, {{math|1=w = 0}} describes a matter-dominated universe, where the pressure is negligible with respect to the mass density. From the generic solution one easily sees that in a matter-dominated universe the scale factor goes as

$$a(t)\; \backslash propto\; t^\{2/3\}\; \backslash qquad\; \backslash text\{matter-dominated\}$$

Another important example is the case of a radiation-dominated universe, namely when {{math|1=w = {{sfrac|1|3}}}}. This leads to

$$a(t)\; \backslash propto\; t^\{1/2\}\; \backslash qquad\; \backslash text\{radiation-dominated\}$$

Note that this solution is not valid for domination of the cosmological constant, which corresponds to an {{math|1=w = −1}}. In this case the energy density is constant and the scale factor grows exponentially.

Solutions for other values of {{mvar|k}} can be found at {{cite web | last=Tersic | first=Balsa | title=Lecture Notes on Astrophysics | url=https://www.academia.edu/5025956|access-date=24 February 2022}}

If the matter is a mixture of two or more non-interacting fluids each with such an equation of state, then

$$\backslash dot\{\backslash rho\}\_\{f\}\; =\; -3\; H\; \backslash left(\; \backslash rho\_\{f\}\; +\; \backslash frac\{p\_\{f\}\}\{c^2\}\; \backslash right)$$

holds separately for each such fluid {{mvar|f}}. In each case,

$$\backslash dot\{\backslash rho\}\_\{f\}\; =\; -3\; H\; \backslash left(\; \backslash rho\_\{f\}\; +\; w\_\{f\}\; \backslash rho\_\{f\}\; \backslash right)\; \backslash ,$$

from which we get

$$\{\backslash rho\}\_\{f\}\; \backslash propto\; a^\{-3\; \backslash left(1\; +\; w\_\{f\}\backslash right)\}\; \backslash ,.$$

For example, one can form a linear combination of such terms

$$\backslash rho\; =\; A\; a^\{-3\}\; +\; B\; a^\{-4\}\; +\; C\; a^0\; \backslash ,$$

where {{mvar|A}} is the density of "dust" (ordinary matter, {{math|1=w = 0}}) when {{math|1=a = 1}}; {{mvar|B}} is the density of radiation ({{math|1=w = {{sfrac|1|3}}}}) when {{math|1=a = 1}}; and {{mvar|C}} is the density of "dark energy" ({{math|1=w = −1}}). One then substitutes this into

$$\backslash left(\backslash frac\{\backslash dot\{a\}\}\{a\}\backslash right)^2\; =\; \backslash frac\{8\; \backslash pi\; G\}\{3\}\; \backslash rho\; -\; \backslash frac\{kc^2\}\{a^2\}$$

and solves for {{mvar|a}} as a function of time.

File:Aleksandr Fridman.png]]

Friedmann published two cosmology papers in the 1922-1923 time frame. He adopted the same homogeneity and isotropy assumptions used by Albert Einstein and by Willem de Sitter in their papers, both published in 1917. Both of the earlier works also assumed the universe was static, eternally unchanging. Einstein postulated an additional term to his equations of general relativity to ensure this stability. In his paper, de Sitter showed that spacetime had curvature even in the absence of matter: the new equations of general relativity implied that a vacuum had properties that altered spacetime.{{Cite book |last1=Tropp |first1=Eduard A. |url=https://www.cambridge.org/core/product/identifier/9780511608131/type/book |title=Alexander A Friedmann: The Man who Made the Universe Expand |last2=Frenkel |first2=Viktor Ya. |last3=Chernin |first3=Artur D. |date=1993-06-03 |publisher=Cambridge University Press |isbn=978-0-521-38470-4 |edition=1 |translator-last=Dron |translator-first=Alexander |doi=10.1017/cbo9780511608131 |translator-last2=Burov |translator-first2=Michael}}{{rp|152}}

The idea of static universe was a fundamental assumption of philosophy and science. However, Friedmann abandoned the idea in his first paper "On the curvature of space". Starting with Einstein's 10 equations of relativity, Friedmann applies the symmetry of an isotropic universe and a simple model for mass-energy density to derive a relationship between that density and the curvature of spacetime. He demonstrates that in addition to one solution is static, many time dependent solutions also exist.{{rp|p=157}}

Friedmann's second paper, "On the possibility of a world with constant negative curvature," published in 1924 explored more complex geometrical ideas. This paper establish the idea that that the finiteness of spacetime was not a property that could be established based on the equations of general relativity alone: both finite and infinite geometries could be used to give solutions. Friedmann used two concepts of a three dimensional sphere as analogy: a trip at constant latitude could return to the starting point or the sphere might have an infinite number of sheets and the trip never repeats.{{rp|p=167}}

Friedmann's paper were largely ignored except – initially – by Einstein who actively dismissed them. However once Edwin Hubble published astronomical evidence that the universe was expanding, Einstein became convinced. Unfortunately for Friedmann, Georges Lemaître discovered some aspects of the same solutions and wrote persuasively about the concept of a universe born from a "primordial atom". Thus historians give these two scientists equal billing for the discovery.{{Cite journal |last=Belenkiy |first=Ari |date=2012-10-01 |title=Alexander Friedmann and the origins of modern cosmology |url=https://pubs.aip.org/physicstoday/article/65/10/38/413748/Alexander-Friedmann-and-the-origins-of-modern |journal=Physics Today |language=en |volume=65 |issue=10 |pages=38–43 |doi=10.1063/PT.3.1750 |bibcode=2012PhT....65j..38B |issn=0031-9228|url-access=subscription }}

Several students at Tsinghua University (CCP leader Xi Jinping's alma mater) participating in the 2022 COVID-19 protests in China carried placards with Friedmann equations scrawled on them, interpreted by some as a play on the words "Free man". Others have interpreted the use of the equations as a call to “open up” China and stop its Zero Covid policy, as the Friedmann equations relate to the expansion, or “opening” of the universe.{{Cite news |last=Murphy |first=Matt |date=November 28, 2022 |title=China's protests: Blank paper becomes the symbol of rare demonstrations |url=https://www.bbc.com/news/world-asia-china-63778871 |work=BBC News}}

• Mathematics of general relativity
• Solutions of the Einstein field equations

{{reflist}}

• {{cite book |first=Dierck-Ekkehard |last=Liebscher |chapter=Expansion |title=Cosmology |location=Berlin |publisher=Springer |year=2005 |isbn=3-540-23261-3 |pages=53–77 |chapter-url=https://books.google.com/books?id=VK_rbBR61eUC&pg=PA53 }}

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