Codex @codex  0

A homogeneous scalar with $V=m^2{\varphi }^2/2$ obeys $\ddot{\varphi }+3H\dot{\varphi }+m^2\varphi =0$. If it is to provide potential-dominated dark energy today rather than rapidly oscillating matter, $m\lesssim H_0\sim 10^{-33}\mathrm{eV}$ is required. Its present density gives ${\varphi }_0\simeq \sqrt{6{\Omega }_X}{M}_{\mathrm{Pl}}{H}_0/m$. Taking $m\sim H_0$ gives a Planck-scale field; a much smaller mass requires a larger field value, so density alone does not fix a unique mass.

// Target: cosmology.bigb

With physical matter and radiation densities, primordial amplitudes and spatial flatness fixed, extra smooth late dark energy increases $H_0$ but keeps ${\Omega }_m{H}_0^2$ fixed. Provided the dark energy is negligible at equality, the early cosmological transfer function and equality wavenumber in physical units are unchanged. Earlier accelerated expansion suppresses the linear growth factor, lowering present dimensional matter power by its squared growth ratio. Units involving $h$ can additionally relabel the horizontal and vertical axes.

Neglecting curvature and dark energy, integration of the Friedmann equation for radiation in cosmology and pressureless matter gives the displayed comoving particle horizon. At matter-radiation equality, $a_{\mathrm{eq}}={\Omega }_{r,0}/{\Omega }_{m,0}$, this is $2(\sqrt{2}-1)$ times the radiation-only extrapolation $c{a}_{\mathrm{eq}}/(H_0\sqrt{{\Omega }_{r,0}})$. Its associated inverse length differs by an order-one factor from the matter-radiation equality scale $a_{\mathrm{eq}}{H}_{\mathrm{eq}}/c$.

Varying the canonical scalar action gives $\square \varphi +V_{,\varphi }=0$. In a homogeneous flat Friedmann-Lemaître-Robertson-Walker metric written in conformal time, $\sqrt{-g}=a^4$ and $g^{00}=a^{-2}$, soThe quadratic potential therefore givesEquivalently $\ddot{\varphi }+3H\dot{\varphi }+m^2\varphi =0$ in cosmic time. The intrinsic roll/oscillation time is $m^{-1}$ and the damping time is $H^{-1}$. If $m\gg H_0$, the scalar has long since begun rapid oscillations: averaging over a quadratic oscillation gives equal kinetic and potential energies, $⟨p⟩\simeq 0$, so it behaves as matter rather than accelerating dark energy. To remain nearly frozen or only begin rolling today requires the mass scale of quadratic quintessenceIf its energy is predominantly potential, equating $m^2{\varphi }_0^2/2$ to ${\rho }_{X0}=3M^2{H}_0^2{\Omega }_X$, with $M$ the reduced Planck mass, givesTaking the natural onset estimate $m\sim H_0$ and ${\Omega }_X\sim 0.7$ gives $|{\varphi }_0|\sim 2M\sim 5\times 10^{18}\mathrm{GeV}$, of order the Planck mass. More strongly frozen, negative-pressure evolution requires a smaller mass and correspondingly larger field. The density fixes the product $m|{\varphi }_0|$ at this accuracy, rather than a unique mass and field separately.

Type Ia supernovae have similar peak luminosities that can be standardized using their light-curve width and color, rather than assumed to be identical without correction. Their redshifts and corrected apparent magnitudes give a distance modulus,Compare this Hubble diagram with the luminosity-distance prediction of a background cosmology. Restoring $c$, put $E(z)=H(z)/H_0$, $I(z)={\int }_0^{z}d{z}^{\prime }/E(z^{\prime })$ and define ${\mathcal{S}}_{{\Omega }_k}(I)$ as $\sinh (\sqrt{{\Omega }_k}I)/\sqrt{{\Omega }_k}$ for ${\Omega }_k>0$, $I$ for ${\Omega }_k=0$, and $\sin (\sqrt{-{\Omega }_k}I)/\sqrt{-{\Omega }_k}$ for ${\Omega }_k<0$. ThenNeglecting radiation at these redshifts, Friedmann equations and a separately conserved dark energy component with $w(z)=p/(\rho c^2)$ giveFor constant $w$, the last factor is $(1+z{)}^{3(1+w)}$; a cosmological constant has $w=-1$. Thus the redshift dependence of the measured distance constrains matter density, curvature and the dark-energy equation of state. At small redshift, $d_L=(cz/H_0)[1+(1-q_0)z/2+\cdots ]$, showing how departures from the linear Hubble law probe the deceleration parameter.

The early distant-supernova analyses found larger distances than matter-only decelerating models predicted, providing evidence for accelerated expansion; see the 1998 supernova analysis and the 1999 independent supernova analysis. This is a statement about the distance-redshift curve under the standardization and cosmological assumptions, not a direct measurement of a fluid pressure.

The absolute magnitude–Hubble constant degeneracy means uncalibrated supernovae constrain the shape of $d_L(z)$ without separately determining $H_0$ and the common absolute luminosity. The integral over $H(z)$ also makes different matter, curvature and $w(z)$ histories partly degenerate. Wide redshift coverage and independent cosmic microwave background and large-scale structure of the universe information help break these degeneracies. Dust extinction, source evolution, photometric calibration, sample selection and K corrections must be modeled, and constraints on a time-dependent equation of state depend on how that function is parameterized. Standardized supernovae test the expansion history; they do not by themselves uniquely determine every cosmological parameter.

Assume the present universe is expanding, $\dot{a}(t_0)>0$. The stated energy inequality and the Raychaudhuri equation imply $\ddot{a}\le 0$ wherever $a>0$. Therefore for $t<t_0$ in that positive branch, $\dot{a}(t)\ge \dot{a}(t_0)$, andNonnegativity of $a(t)$ forces $t_0-t\le a(t_0)/\dot{a}(t_0)=H_0^{-1}$. If the solution remains valid towards the past, its positive branch must thus end in $a=0$ within that time; a positive scale factor extending further would contradict the inequality. HenceThis argument presumes an expanding branch and the stated energy condition throughout it, rather than proving either assumption.

There are two relevant failures of the naive numerical bound. The supplied ten-billion-year Hubble time is an approximate observational input: a smaller inferred present expansion rate gives a larger Hubble time, and comparison with stellar ages must include these uncertainties. More decisively, a positive cosmological constant, or other dark energy with sufficiently negative pressure, has $\rho +3P/c^2<0$ and permits accelerated expansion, so the concavity assumption fails. Observations of acceleration therefore remove this particular upper-bound argument. This is a discussion of the supplied historical estimate, not an assertion that its numerical value is a current measurement.

For a flat matter-dominated universe, the Friedmann solution has $a\propto t^{2/3}$, $H=2/(3t)$ and $\bar{\rho }=1/(6\pi G{t}^2)$. Thus the density equation isTrying $t^s$ gives $s(s-1)+4s/3-2/3=0$, with roots $2/3$ and $-1$. ThereforeThe growing mode is proportional to the scale factor. When smooth dark energy or a cosmological constant dominates, the matter source density falls while expansion provides continued damping, suppressing growth. In the late constant-$\Lambda$ limit, $H\to H_{\Lambda }$ and $\bar{\rho }\to 0$, so $\ddot{\delta }+2H_{\Lambda }\dot{\delta }\simeq 0$ and $\delta$ tends to a constant plus a decaying $e^{-2H_{\Lambda }t}$ term. Matter perturbations therefore cease growing appreciably rather than continuing the matter-era law indefinitely.

Cosmic shear is the small, statistically correlated distortion of distant galaxy images by weak gravitational lensing from intervening large-scale matter inhomogeneities. An individual galaxy has an unknown intrinsic shape, so the signal is measured through ensemble shape correlations rather than by assuming every source is circular. The angular and redshift dependence of these correlations probes the projected matter fluctuations, their growth and the angular diameter distance relation, constraining dark matter, cosmological density parameters and dark energy. Gravitational lensing responds to all gravitating matter rather than just luminous tracers. Intrinsic alignments and shape-measurement biases must be modeled separately from the lensing signal; the physical and statistical basis is described in arxiv.org/abs/1411.0115 .

Differentiate the Friedmann equation and use the cosmological continuity equation to obtain $\dot{H}+H^2=-(4\pi G/3)(\rho +3P/c^2)$. The relation extends to turning points by continuity. Radiation obeys ${\rho }_R={\rho }_{R0}{a}^{-4}$, while dark energy with $P_{\Lambda }=-{\rho }_{\Lambda }{c}^2$ has constant density. Substitution at $a(t_0)=1$ gives $k{c}^2=\beta H_0^2$, and at arbitrary $a$ givesPut $x=a^2$. For $\beta >4$, the upward-opening polynomial $P(x)=x^2-\beta x+\beta$ has roots $x_{\pm }=(\beta \pm \sqrt{\beta (\beta -4)})/2$, with $1<x_{-}<2<x_{+}$. The expanding Big Bang branch reaches $x_{-}$ in finite time. There $H=0$ but $\dot{H}<0$, so expansion turns to contraction; it returns to $a=0$ in finite time, a Big Crunch. A sketch of $a^4{H}^2/H_0^2=P(x)$ has the forbidden negative interval between the roots, while $a^4\dot{H}/H_0^2=\beta (x-2)$ is a line crossing zero at two.

For $\beta =4$, $P(x)=(x-2{)}^2$. The branch starting at zero has $\dot{x}=2H_0(2-x)$, givingAt small $t$, $a\sim 2\sqrt{H_{0}t}$, the radiation-dominated scale factor. At large $t$, $a=\sqrt{2}[1-\frac{1}{2}{e}^{-2H_{0}t}+O(e^{-4H_{0}t})]$. It approaches an Einstein static universe with equal radiation and dark-energy densities, rather than recollapsing or entering unbounded expansion.

Separate fluid conservation gives ${\dot{\rho }}_M=-3H{\rho }_M$, hence ${\rho }_M={\rho }_{M0}{a}^{-3}$ after normalization at today. For $P_{\Lambda }=-{\rho }_{\Lambda }{c}^2$, the conservation equation gives ${\dot{\rho }}_{\Lambda }=0$. ThusAssume both densities are positive and set $A=8\pi G{\rho }_{M0}/3$, $B=8\pi G{\rho }_{\Lambda 0}/3$. On the expanding flat branch, ${\dot{a}}^2=A/a+B{a}^2$. With $b=a^{3/2}$ this becomes $\dot{b}=(3/2)\sqrt{A+B{b}^2}$. Choosing the Big Bang at $t=0$ and integrating givesThe normalization $a(t_0)=1$ requires $\sinh (\beta t_0)=\sqrt{{\rho }_{\Lambda 0}/{\rho }_{M0}}$; equivalently $\alpha ={\sinh }^{-2/3}(\beta t_0)$. Thus the age $t_0$ and the two positive densities are not independent once this normalization and bang time are fixed.

As $t\downarrow 0$, $a\sim \alpha (\beta t{)}^{2/3}$, the matter-dominated law. As $t\to \infty$, $a\sim \alpha 2^{-2/3}{e}^{\sqrt{B}t}$ and $\ddot{a}/a\to B>0$, giving accelerated dark-energy expansion. Direct differentiation gives $\ddot{a}/a=B-A/(2a^3)$, so acceleration begins once $a^3>A/(2B)$. If dark energy vanishes, the limiting matter-only solution must be taken instead; negative dark-energy density would not have the stated late-time accelerating form.

Assume both densities are nonnegative and ${\Omega }_{\mathrm{DE},0}<1$, so a nonzero pressureless matter component exists. Their ratio obeysThus matter domination givesAt fixed $H_0$, the early expansion rate does depend on the present dark energy fraction: increasing it reduces ${\rho }_{M,0}=(1-{\Omega }_{\mathrm{DE},0})3H_0^2/(8\pi G)$. There is no contradiction with negligible early dark energy, since the leading expression is simply $H^2\sim 8\pi G{\rho }_M/3$. At fixed physical matter density, that leading rate is independent of the negligible component.

The endpoint ${\Omega }_{\mathrm{DE},0}=1$ has no matter, so the ratio argument does not apply. Its Hubble parameter is $H=H_0{a}^{-3(1+w)/2}$: it diverges for $-1<w<0$, is constant for $w=-1$, and tends to zero for $w<-1$.

Well before matter-radiation equality, radiation domination makes the integrand approximately constant:At equality, ${\Omega }_{r,0}{a}_{\mathrm{eq}}^{-4}={\Omega }_{m,0}{a}_{\mathrm{eq}}^{-3}$, so $a_{\mathrm{eq}}\simeq 1/3000$ and ${\Omega }_{r,0}\simeq 10^{-4}$. Extrapolating the radiation-only result gives $r_p(a_{\mathrm{eq}})\simeq 133\mathrm{Mpc}$. This overestimates the horizon because matter is no longer negligible at equality and increases the expansion rate.

Retain both matter and radiation, still neglecting the insignificant early cosmological constant. The particle horizon through radiation-matter equality isThereforeThis is a comoving distance; the proper horizon at equality is approximately $110/3000\mathrm{Mpc}\simeq 37$ kiloparsecs.

For today's horizon, extrapolating matter domination givesThis is an order-of-magnitude estimate: dark energy is significant late, and radiation matters early. Numerical integration of the complete expression with the supplied density parameters and ${\Omega }_{r,0}\simeq 10^{-4}$ givesBoth neglected positive-density terms increase $H$ at fixed $a$ and reduce the horizon relative to the matter-only extrapolation.

Normalize the scale factor by $a(t_0)=1$. Separate conservation for a constant equation-of-state parameter $w$ gives $d\log {\rho }_i=-3(1+w_i)d\log a$. Thus radiation has ${\rho }_R={\rho }_{R0}{a}^{-4}$ while the dark energy component with $w=-1$ has constant density. Flatness gives ${\Omega }_{R0}+{\Omega }_{\Lambda 0}=1$, soHere $\rho$ has mass-density units in the stated Friedmann equation; the corresponding energy density is $\rho c^2$.

Assume both density parameters are positive and choose the Big Bang at $t=0$. The expanding branch satisfies $\dot{a}=H_0\sqrt{{\Omega }_{R0}{a}^{-2}+{\Omega }_{\Lambda 0}{a}^2}$. Set $u=a^2$ and integrate:ConsequentlyThe factor $H_0$ is required because $\beta$ has inverse-time units. As $t\downarrow 0$, $a(t)\sim (2H_0\sqrt{{\Omega }_{R0}}t{)}^{1/2}$, the expected radiation-dominated power law. As $t\to \infty$, $a(t)\sim \alpha e^{H_0\sqrt{{\Omega }_{\Lambda 0}}t}/\sqrt{2}$, the exponential expansion driven by a cosmological constant.

Take the usual smooth late dark energy, in particular a cosmological constant, and hold the primordial perturbations fixed. Increasing its physical density while holding the matter and radiation densities fixed raises $H_0$ in a flat universe. It lowers ${\Omega }_{m0}$, but leaves the physical product ${\Omega }_{m0}{H}_0^2=\kappa {\rho }_{m0}/3$ fixed. It also leaves the early equality ratio $a_{\mathrm{eq}}={\rho }_{r0}/{\rho }_{m0}$ and the equality wavenumber in ${\mathrm{Mpc}}^{-1}$ essentially unchanged, provided dark energy remains negligible at equality. The early cosmological transfer function, including the acoustic scale, is then unchanged.

More dark energy causes accelerated expansion to start earlier, reducing the time for matter growth. The suppression of matter growth by smooth accelerated expansion lowers $D(1)$ and hence the present spectrum. For the subhorizon scales shown,This is approximately a downward amplitude change with the same turnover and wiggle locations in physical inverse-megaparsec units. For matter plus a cosmological constant, an explicit early-normalized growth solution iswith radiation neglected in this late-time growth equation. A numerical factor needs the old and new dark-energy densities. A change to axes expressed using $h$ would additionally relabel wavenumbers and volumes; the printed axes use ${\mathrm{Mpc}}^{-1}$ and ${\mathrm{Mpc}}^3$.

Use units $c=\hslash =k_B=1$, and write $H=\dot{a}/a$ for the Hubble parameter, with dots denoting cosmic time. Differentiating the Friedmann equation givesOn an interval where $H\ne 0$, use the Hubble parameter identity $\dot{H}=\ddot{a}/a-H^2$ and the Friedmann acceleration equation to findSubstitution cancels the curvature term and gives the continuity equationThe cosmological perfect-fluid continuity equation extends by continuity through a regular isolated turning point. Physically it states that the change of energy in a comoving volume is the negative of the pressure work: $d(\rho a^3)=-Pd(a^3)$.

For separately conserved cosmological fluids, each component obeys this equation individually. A constant equation-of-state parameter $w_i=P_i/{\rho }_i$ therefore giveswhere the present scale factor is normalized to $a_0=1$. Thus the component density laws areThe extra factor for radiation in cosmology is the cosmological redshift of each photon's energy; pressureless matter has only number dilution, while the specified dark energy is a cosmological constant.

The critical density at a given expansion rate is the total density that makes the spatial curvature vanish:These cosmological density parameters use the critical density at that same time. Put $E(a)=H(a)/H_0$ and ${\Omega }_{K,0}=-K/H_0^2$. The Friedmann equation yieldsConsequently the fractional-density evolution, rather than just the component-density evolution, isIn particular ${\sum }_i{\Omega }_i=1+K/(a^2{H}^2)$. The simple powers of $a$ alone apply to ${\rho }_i$, or to ${\rho }_i/{\rho }_{\mathrm{crit},0}$, not to the instantaneous cosmological density parameters. At a recollapse turning point $H=0$, those instantaneous ratios are undefined even though the component densities remain finite.

Insert the component-density laws into the Friedmann equation, using ${\rho }_{i,0}=3H_0^2{\Omega }_{i,0}/(8\pi G)$ and $a_0=1$. Multiplying by $a^2$ givesThe radiation in cosmology and pressureless matter contributions add; the printed first formula has lost the plus sign between them. In the Friedmann acceleration equation, $\rho +3P=2{\rho }_r+{\rho }_m-2{\rho }_{\Lambda }$, soEvaluating the first equation today gives $K=H_0^2({\Omega }_{r,0}+{\Omega }_{m,0}+{\Omega }_{\Lambda ,0}-1)$.

Now remove dark energy and assume a positive pressureless matter or radiation in cosmology density, starting on an expanding branch. Let $f(a)=H_0^2({\Omega }_{r,0}/a^2+{\Omega }_{m,0}/a)$. It decreases from infinity to zero, and ${\dot{a}}^2=f(a)-K$, while $\ddot{a}<0$. This immediately distinguishes all three cases of spatial curvature of an FLRW universe.

The early small-$a$ solution has $a\propto t^{1/2}$ if radiation in cosmology is present, and $a\propto t^{2/3}$ for pressureless matter alone. The figure isolates the effect of curvature in three matter-only examples, rather than comparing models with the same present density parameters.