Barotropic equation of state

A constant barotropic equation of state has $P=w\rho$. The cases $w=0$ and $w=1/3$ describe pressureless matter and radiation respectively.

A coasting fluid has pressure $p=-\rho /3$. The cosmological perfect-fluid continuity equation makes its density scale as $a^{-2}$. If it alone fills a flat expanding universe, the scale factor grows linearly in physical time. Its background density has the same scale-factor dependence as a spatial-curvature term, although the fluid does not itself create spatial curvature.

For conserved $w=-1/3$ scalar energy, ${\rho }_X=3M^2{H}_0^2{\Omega }_X/a^2$ and ${\varphi }^{\prime }=\sqrt{2}M{H}_0\sqrt{{\Omega }_X}$ on the increasing branch. With dust and radiation, define $b={\Omega }_m/(2{\Omega }_X)$ and $A=\sqrt{{\Omega }_m^2/4-{\Omega }_r{\Omega }_X}/{\Omega }_X$. When the discriminant is positive,and $V(\varphi )=2M^2{H}_0^2{\Omega }_X/a(\varphi {)}^2$. The formulas retain all density parameters without eliminating one by the flatness sum.

An isotropic, noninteracting network of domain walls with energy density proportional to $a^{-1}$ has effective equation-of-state parameter $w=-2/3$. This follows from constant-equation-of-state density scaling $\rho \propto a^{-3(1+w)}$.

For constant $w\ne -1$, the cosmological continuity equation integrates towhen $a=1$ at the reference time and ${\rho }_0$ is the reference density.

For an expanding spatially flat universe with constant $w>-1$ and $a(t_0)=1$,It vanishes at the finite past time $t_{\ast }=t_0-2/[3(1+w)H_0]$.

Pressureless matter has equation-of-state parameter $w=0$, energy density $\rho \propto a^{-3}$ and negligible stress compared with its rest-mass density.

An isotropic network of long cosmic strings has average equation-of-state parameter $w=-1/3$ and energy density proportional to $a^{-2}$.

The equation of state averages the stress of a cosmic string over an isotropic network; it is not the definition of the individual defect.

An isotropic cosmic string network with $\rho =C/a^2$ has equation-of-state parameter $w=-1/3$. The Friedmann equation gives ${\dot{a}}^2=8\pi GC/3-K$. On the expanding branch with positive right side, $a=B(t-t_B)$, the comoving Hubble radius is constant, and the cosmological density parameter is constant. There is no growing flatness problem, but flatness is not an attractor. If this exact coasting solution extends all the way to $t_B$, then ${\int }_{t_B}^{t}d{t}^{\prime }/a(t^{\prime })$ diverges and there is no finite particle horizon. A finite starting time changes that global conclusion.

A closed universe containing only a cosmic string network can remain static. Its Friedmann acceleration equation vanishes identically because $P_s=-{\rho }_s/3$, and its Friedmann equation requires ${\rho }_s{a}^2=3K/(8\pi G)$. Once this conserved normalization is tuned, every positive radius gives a static solution. The usual Hubble-normalized cosmological density parameter is undefined at a static epoch.

At fixed string normalization $C={\rho }_s{a}^2$, scale-factor shifts move between static solutions and are neutral, with no restoring force. The acceleration equation gives $\ddot{a}=0$, but the Friedmann constraint also imposes ${\dot{a}}^2=8\pi GC/3-K$. At the static normalization it excludes a nonzero velocity. A slightly larger $C$ admits linearly expanding or contracting branches; a smaller $C$ admits no real closed solution. This distinguishes permitted perturbations from constraint violations.

Consider a spatially flat universe containing radiation and a cosmic string network. If their densities are both ${\rho }_{\mathrm{eq}}$ at scale factor $a_{\mathrm{eq}}$, thenIn conformal time $d\tau =dt/a$, the flat Friedmann equation becomesChoosing the Big Bang at $\tau =0$, its expanding solution is