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Besides them, we plot the histograms of the RM distribution of the related source. | Besides them, we plot the histograms of the RM distribution of the related source. |
Table 7 lists the mean «RM» value, opm, and the maximum absolute value of the RM distributions. | Table \ref{RMvalues} lists the mean $<$ $>$ value, $\sigma_{\rm RM}$, and the maximum absolute value of the RM distributions. |
The radio galaxies and the filaments have similar «RM» and egy values (<RM>~+20 rad m7, orm~10 rad m?). | The radio galaxies and the filaments have similar $<$ $>$ and $\sigma_{\rm RM}$ values$<$ $> \sim +20$ rad $^{-2}$, $\sigma_{\rm RM}\sim10$ rad $^{-2}$ ). |
These values are typical of radio sources | These values are typical of radio sources |
have moceled as a black body. which is far too much to be attributable to the host galaxy or to scatteing. | have modeled as a black body, which is far too much to be attributable to the host galaxy or to scatteing. |
The fit to the X-ray spectrum is therefore typical of a Sevlert 1 galaxy. | The fit to the X-ray spectrum is therefore typical of a Seyfert 1 galaxy. |
There is an underlying power law. to which a soft excess. probably the direct. quasi-thermal radiation from the aceretion disk. has to be added. | There is an underlying power law, to which a soft excess, probably the direct quasi-thermal radiation from the accretion disk, has to be added. |
The Fe Ix emission line complex is rather broad. (ELINM. ~οὐ) ο). as it is often. found. in Sevfert 2 galaxies. lor example in the prototypical NGC LOGS (às à result of the superposition of various components. Iwasawa et al 1997). | The Fe K emission line complex is rather broad (FHWM $\sim 900\, {\rm eV}$ ), as it is often found in Seyfert 2 galaxies, for example in the prototypical NGC 1068 (as a result of the superposition of various components, Iwasawa et al 1997). |
Phe equivalent width of the complex is 380EX359OW which lies in between the one expected from a Sevfert 1 and a Compton-thin Sevíert 2. all consistent with the Sevfert 1.8/1.9 nature of the source. | The equivalent width of the complex is $380^{+330}_{-320}\, {\rm eV}$ which lies in between the one expected from a Seyfert 1 and a Compton-thin Seyfert 2, all consistent with the Seyfert 1.8/1.9 nature of the source. |
In order to gain further insight into the N-rav. properties of this X-ray source. we have extracted its EPIC pn ligh curve. | In order to gain further insight into the X-ray properties of this X-ray source, we have extracted its EPIC pn light curve. |
Time intervals have been binned in 300 s. resulting in a signal to noise of Ls. | Time intervals have been binned in 300 s, resulting in a signal to noise of $\sim 18$. |
The background counts have been estimated from the same background subtraction region tha was used to analyze the spectrum. | The background counts have been estimated from the same background subtraction region that was used to analyze the spectrum. |
The resulting 0.2-12 keV light curve is shown in fig 9.. | The resulting 0.2-12 keV light curve is shown in fig \ref{fig-lightcurve}. |
The average count rate is 1.67 ct/s. However. the ligh curve is not consistent with a constant Lux. | The average count rate is 1.67 ct/s. However, the light curve is not consistent with a constant flux. |
A first elance at Πο 9 reveals that 2/3 of the properly computed La error bars do not cross the horizontal fit. while only 1/3 woul be expected. | A first glance at fig \ref{fig-lightcurve} reveals that 2/3 of the properly computed $\sigma$ error bars do not cross the horizontal fit, while only $1/3$ would be expected. |
‘Phe 47. of the fit to a constant. intensity. is extremely poor. Afv=341/65. | The $\chi^2$ of the fit to a constant intensity is extremely poor, $\chi^2/\nu = 341/65$. |
To estimate the rms variability. we compare the measured variance of the count rates to the variance expected. from the error. bars ancl find 5 per cent. intrinsic variability. on that time. scale. | To estimate the rms variability, we compare the measured variance of the count rates to the variance expected from the error bars and find 5 per cent intrinsic variability on that time scale. |
Although small. these variations are highly significant. | Although small, these variations are highly significant. |
Note that for a 107M. black hole. variability is expected down to scales of seconds. | Note that for a $10^5\, {\rm
M}_{\odot}$ black hole, variability is expected down to scales of seconds. |
The fact that the source varies. brings additional support to the rejection of a Compton thick model. which would not predict short term: variability. | The fact that the source varies brings additional support to the rejection of a Compton thick model, which would not predict short term variability. |
Reconciling the Sevlert 1.8/1.9 character of 111320|551. as drawn from the optical observation in 1998. with the SALAI- data. taken late 2001. has two possible scenarios. | Reconciling the Seyfert 1.8/1.9 character of H1320+551, as drawn from the optical observation in 1998, with the XMM-Newton data, taken late 2001, has two possible scenarios. |
exception is the Botwensetal.(2001b) data which 1ecessitated an extinction correction of 0.5 dex in order to briD>oO it onto the same system as the other data. | exception is the \citet{bouwens7} data which necessitated an extinction correction of 0.8 dex in order to bring it onto the same system as the other data. |
Where appropriate. we have also applied incompeleness corrections assuming a 5ciechter extrapolation with a=—1.5.For example. Bouwenseal.(200Lh) ouly compute the UV luminosity fiction down to £=0.3L,. | Where appropriate, we have also applied incompleteness corrections assuming a Schechter extrapolation with $\alpha=-1.5$.For example, \citet{bouwens7} only compute the UV luminosity function down to $L=0.3L_\star$. |
Ext‘apolaing over tie. [ull rauge of £L implies a correction of 0.36 dex. | Extrapolating over the full range of $L$ implies a correction of 0.36 dex. |
The tpper pane shows our leastirelents ol the &obal inass deusity as solid points. together with a nuuber of measurements from the literature (Rucduicketal.2003:Dickinson)Foutanaeal.20Ol:Droryet2001.2005) as assoled open points. | The upper panel shows our measurements of the global mass density as solid points, together with a number of measurements from the literature \citep{rudnick,dickmass,k20,drory2004,drory2005} as assorted open points. |
The authors of the above ΑΟ)ss have corrected their data for the unobserved portion of the mass function. | The authors of the above works have corrected their data for the unobserved portion of the mass function. |
In their analysis of tje Cetini Deep Deep Survey. Glazebrooketal.(2001) have chosen 1ot to extrapolate beyond what they observe. | In their analysis of the Gemini Deep Deep Survey, \citet{gddsmass} have chosen not to extrapolate beyond what they observe. |
The GDDS data lor galaxies more massive than log(Al/Al.)>10.2 are plotted as lower linits. | The GDDS data for galaxies more massive than $\log(M/M_\odot) >10.2$ are plotted as lower limits. |
The corrections for the choice of IME. diseussecl in sectio 3.2.. have been applied. | The corrections for the choice of IMF, discussed in section \ref{ssec:masserr}, have been applied. |
The solid line in both panels of Figure 3) come lrom Hartwick(2001).. who derived the val star [formation history from observations of the local universe. | The solid line in both panels of Figure \ref{fig:f3} come from \citet{hartwick2004}, who derived the global star formation history from observations of the local universe. |
Bjelly. he moclel uses the ributiou in metalliciy ol stars to derive dAL/dZ(where Z is the yetalliciy) and the relationship or globular clusters to derive dZ/dl (wlere / is the age of the universe). | Briefly, the model uses the distribution in metallicity of stars to derive $dM_\star/dZ$(where $Z$ is the metallicity) and the age-metallicity relationship for globular clusters to derive $dZ/dt$ (where $t$ is the age of the universe). |
ubining Λι/dZ and dZ/dl. oue obtains dll./d!= SER. the star formation ‘ate. | Combining $dM_\star/dZ$ and $dZ/dt$, one obtains $dM_\star/dt \equiv$ SFR, the star formation rate. |
This siuiple -jodel does a ve“Vv οοος job of explaining the star fortlation |istory o‘the ulVerse. as sLOWLL —Nw the agreemen between it and the observatious in the lower j»auel. | This simple model does a very good job of explaining the star formation history of the universe, as shown by the agreement between it and the observations in the lower panel. |
UsingOm |this star [ormation Ιistory as an inp ito the PECGASE 2.0 softwafen we compute a uodel global selar nass clesity. | Using this star formation history as an input to the PEGASE 2.0 software, we compute a model global stellar mass density. |
The result is ploted in the upper panel. | The result is plotted in the upper panel. |
It is iu excellei| agreeuent with our οobal stellar lass ensity measurelrents. | It is in excellent agreement with our global stellar mass density measurements. |
Note that this agreement is not «lepeucdeit on the details of the Harwick model. | Note that this agreement is not dependent on the details of the Hartwick model. |
Almost any escription of the sta: fortation history wlich agrees with he observe star formation rates will. «once integrated. produce exXxd agreement wil &lobal stellar 1jass cdeusiv. | Almost any description of the star formation history which agrees with the observed star formation rates will, once integrated, produce good agreement with global stellar mass density. |
This is illustrated by the «ottedl liue in Figure :Y. | This is illustrated by the dotted line in Figure \ref{fig:f3}. |
On the lower panel. us shows a crude. thiree-segiment ~conmect-the-dots” description of the sta ‘fortvation history. | On the lower panel, this shows a crude, three-segment “connect-the-dots” description of the star formation history. |
hi he upper panel the cloted line shows the results ol integrating (again wil1 PEGASE 2.0) this sta ‘formation histo‘vy. | In the upper panel the dotted line shows the results of integrating (again with PEGASE 2.0) this star formation history. |
Again. there is good agreement wit1 our global stellar uiass density calctlatious. | Again, there is good agreement with our global stellar mass density calculations. |
Noe that couverting a star lornation rate to GSALD i sakürlv robust. procecure. | Note that converting a star formation rate to GSMD is a fairly robust procedure. |
It is fairly insensitive to the details of the population imodelliug. | It is fairly insensitive to the details of the population modelling. |
bludeed. there is only a slight oss Of accuracy even if one makes the exlrenie assttuptiou that all stars. once created. never cie. | Indeed, there is only a slight loss of accuracy even if one makes the extreme assumption that all stars, once created, never die. |
This is )ecatise (for all ‘easolalble initial mass [ucttous) the bulk of the mass in stars comes [rom sars of less than l1 M.. which have a miain-sequence iletime of approximately a Hubble time. | This is because (for all reasonable initial mass functions) the bulk of the mass in stars comes from stars of less than 1 $M_\odot$ , which have a main-sequence lifetime of approximately a Hubble time. |
In this simplilied case oue can do a straight integration of the αι/d! (with AL,—0 at /= 0) to determiie. AZ(). | In this simplified case one can do a straight integration of the $dM_\star/dt$ (with $M_\star=0$ at $t=0$ ) to determine $M(t)$ . |
There are two things tonote in the top pauel of Figure 3.. | There are two things tonote in the top panel of Figure \ref{fig:f3}. |
The first is the excellent agreement | The first is the excellent agreement |
hat the eas disks they observe at 2=1 aud 2= have considerably iore molecular eas relative to the stars. vpically about30-50%... compared to ~15% for the A\Gilky Wavy aud nearby disk galaxies (Helferetal.2003). | that the gas disks they observe at $z=1$ and $z=2$ have considerably more molecular gas relative to the stars, typically about, compared to $\sim 1- 5\%$ for the Milky Way and nearby disk galaxies \citep{Helfer2003}. |
. This trend is consistent with our estimates. | This trend is consistent with our estimates. |
Although here is some uncertaintv im the Πο masses of Taccouietal.(2010) because of the uncertainty in the value of New. here appears to be little doubt that the ratio of ID;2 to stellar mass iu the Όσ]ν and DX/DM galaxies is higher han typical values for similar galaxies at +=0. | Although there is some uncertainty in the 2 masses of \cite{t2010} because of the uncertainty in the value of $_{CO}$, there appears to be little doubt that the ratio of 2 to stellar mass in the BzK and BX/BM galaxies is higher than typical values for similar galaxies at $z=0$. |
Duc ο the scusitivity of ρµ5 to the form of MGDR. future observations of pj,2 or the ratio at higher redshift would allow us to better constrain the evolution of AIGDR. aud reduce the area between the bounding curves of Fie. 5.. | Due to the sensitivity of 2 to the form of MGDR, future observations of 2 or the ratio at higher redshift would allow us to better constrain the evolution of MGDR, and reduce the area between the bounding curves of Fig. \ref{OBfig}. |
Iu our open box model. the rYOeservoldr. ο. is augmented by an inflow of gas from at a rate of LOO10537,AlpeP? +. depending ou z. | In our open box model, the reservoir, $\mrhi$, is augmented by an inflow of gas from at a rate of $10^7 - 10^8 M_\sun \permpccu$ $^{-1}$, depending on $z$. |
This ligh rate of inflow means that the eas being accreted is mostly ionized since the fraction of neutral hydrogen outside of the reservoir is too stall. | This high rate of inflow means that the gas being accreted is mostly ionized since the fraction of neutral hydrogen outside of the reservoir is too small. |
The interred from observations of DLA svstcus accounts for the aassociated with ealactic disks. | The inferred from observations of DLA systems accounts for the associated with galactic disks. |
As mentioned in refopenbox.. O'Mearaetal.(2007) find that systems with: σον density. Xj;«κ24n21029 απ2? account for.: 5ο15% ofneutral hydrogen atomis at all +<6. | As mentioned in \\ref{openbox}, , \cite{OMeara2007}
find that systems with column density $\Sigma_{HI} < 2 \times 10^{20}$ $^{-2}$ account for $\approx 15\%$ ofneutral hydrogen atoms at all $z < 6$. |
Therefore. the fraction of ooutside DLA systems is about 1554. corresponding to roughly 1.5«10* M.Mpe.7, | Therefore, the fraction of outside DLA systems is about $15\%$, corresponding to roughly $1.5 \times 10^7$ $_\sun
\permpccu$. |
For an average inflow rate of a few times 10* AL.MpeP2 | for the past 10 Cyr. this intergalactic ccould only account for ten percent of the total. at most. | For an average inflow rate of a few times $10^7$ $_\sun
\permpccu$ $^{-1}$ for the past 10 Gyr, this intergalactic could only account for ten percent of the total, at most. |
Therefore. the inflow of eas needed. for fucling ongoing star formation represented by dpi;/df must be alinost completely iouized. | Therefore, the inflow of gas needed for fueling ongoing star formation represented by $d\mrext/dt$ must be almost completely ionized. |
Receutly. cold. flows have been suggested as an Huportant source of eas for ealaxy formation and evolution (Ixeresetal.2005:Dekel&Biruboim2006). | Recently, cold flows have been suggested as an important source of gas for galaxy formation and evolution \citep{Keres2005,dekel06}. |
. In these models. galactic disks in halos with M<10123. are built up by direct accretion of cold. gaseous streams from the cosimic web. | In these models, galactic disks in halos with $M \la 10^{12} M_\sun$ are built up by direct accretion of cold gaseous streams from the cosmic web. |
For galaxies with larger masses. cold flows are also the dominant meas of nass accretion. but different outcomes for individual galaxies depoeud ou the epoch of inflow. | For galaxies with larger masses, cold flows are also the dominant means of mass accretion, but different outcomes for individual galaxies depend on the epoch of inflow. |
If this picture is correct. our work implies that the cold flows iust be alinost eutirely ionized. | If this picture is correct, our work implies that the cold flows must be almost entirely ionized. |
The same is true if the eas needed to fuel star formation is brought in primarily through minor mergers. | The same is true if the gas needed to fuel star formation is brought in primarily through minor mergers. |
Ifthis eas were in atomic form. it would be part of the inventory of atomic gas observed in the DLA systems. which we have shown ius refopenbox to contribute negligibly to fuclue the star formation at all redshifts up to :—f. | If this gas were in atomic form, it would be part of the inventory of atomic gas observed in the DLA systems, which we have shown in \\ref{openbox} to contribute negligibly to fueling the star formation at all redshifts up to $z = 4$. |
The open box model requires dpi,dt~SPRD. or about 10* to LOS AT,ApeP 1; | The open box model requires $d\mrext/dt \sim SFRD$, or about $10^7$ to $10^8$ $_\sun
\permpccu$ $^{-1}$. |
We use these unubers for dp.4/df to calculate a cooling time for the gas in the contest of two models for gas accretion onto ealaxics: two-phase cooling of hot halo gas (Maller&Bullock and cold dow accretion (I&eres&ITeruquist2009). | We use these numbers for $d\mrext/dt$ to calculate a cooling time for the gas in the context of two models for gas accretion onto galaxies: two-phase cooling of hot halo gas \citep{MB2004} and cold flow accretion \citep{KH2009}. |
. We estimate the cooling time. f,,,5. bv taking where Pyas is the average mass deusitv of the cooling ionized eas sinoothed over the appropriate volume (to be specified for cach cooling model individually). | We estimate the cooling time, $t_{cool}$, by taking where $\rho_{gas}$ is the average mass density of the cooling ionized gas smoothed over the appropriate volume (to be specified for each cooling model individually). |
We set ρα equal to myn.f where a. is the local umber density of electrous aud f£ is the filline factor for the relevant volumes (o0). | We set $\rho_{gas}$ equal to $m_{H} n_e f$ where $n_e$ is the local number density of electrons and $f$ is the filling factor for the relevant volumes $\bar{n}_e/n_e$ ). |
Combining this with the cooling time where ACD)is the cooling fuuctiou of the eas. eives As a basis for comparison. we first estimate the filliue factor f for hot halos of L. galaxies out to the virial radius. | Combining this with the cooling time where $\Lambda(T)$is the cooling function of the gas, gives As a basis for comparison, we first estimate the filling factor $f$ for hot halos of $L_*$ galaxies out to the virial radius. |
We iake the simplistic assumption that the uuivorse is made up of £, galaxies with masses Mi,~LOY XL.aud cieula velocities of ~160 an s | We make the simplistic assumption that the universe is made up of $L_*$ galaxies with masses $M_{dyn} \sim 10^{12}$ $_\sun$and circular velocities of $\sim160$ km $^{-1}$. |
Therefore. the virial radius. Rei,=GM,οἱον300 kpc. | Therefore, the virial radius, $R_{vir} = GM_{dyn}/v^2 \sim 300$ kpc. |
We estimate the average nuuber density in this simple universe composed of £. galaxies by divicing the total Imminosity density, £2. by £.. | We estimate the average number density in this simple universe composed of $L_*$ galaxies by dividing the total luminosity density, $\mathcal{L}$, by $L_*$. |
We use the r baud values frou Blautonctal.(2003) which caleulates the ealaxy hunuinositv function at 2~0.1 frou, SDSS data: £mLSEdO hM ?. fee12.10 2 AL. | We use the r band values from \cite{Blanton2003}
which calculates the galaxy luminosity function at $z\sim0.1$ from SDSS data: $\mathcal{L} \approx 1.84 \times 10^8$ h $_{\odot}$ $^{-3}$, $L_* \approx 1.2 \times 10^{10}$ $^{-2}$ $_{\odot}$. |
This viekls vp.~£/L.0.015 l? ?—5«10? 7. | This yields $n_{L_*} \sim \mathcal{L} / L_* \sim 0.015$ $^3$ $^{-3} \sim 5 \times 10^{-3}$ $^{-3}$. |
More recent work ou the huuinositv function using SDSS DRG vields simular results (Moutero-Dorta 2009). | More recent work on the luminosity function using SDSS DR6 yields similar results \citep{MD2009}. |
.. Therefore. in this simple universe. the filline factor for the L. ealaxy halos is Maller&Bullock(2001) consider eas within the cooling radius of a halo. Πρ. cooling via cloud fracinentation. | Therefore, in this simple universe, the filling factor for the $L_*$ galaxy halos is \cite{MB2004} consider gas within the cooling radius of a halo, $R_c$, cooling via cloud fragmentation. |
This results in the formation of wari. (~104 KJ clouds within the hot eas halo. | This results in the formation of warm $\sim 10^4$ K) clouds within the hot gas halo. |
Iu this model for the two-phase cooling of the hot halo eas. the relevant temperature for the eas is the virial temperature of the halo. ~109 I< for a Milkv Wav type halo. | In this model for the two-phase cooling of the hot halo gas, the relevant temperature for the gas is the virial temperature of the halo, $\sim 10^6$ K for a Milky Way type halo. |
ας&Ieruquist(2009) find that the majority of cold clouds that form around Milkv. Way type galaxies aro the result of filamentary “cold inode accretion. | \cite{KH2009} find that the majority of cold clouds that form around Milky Way type galaxies are the result of filamentary "cold mode" accretion. |
Most of the gas does not reach the virial temperature of the halo. ~109 K. but rather cools from a 1iaxiunun emperature of ~LOL10? K. We examine the possible values for. aud f in these wo models by calculating f as a function of », for three emperatures: LO! IS and 10° Ik for cold flow accretion and 109 K for cooling from the lot halo. | Most of the gas does not reach the virial temperature of the halo, $\sim 10^6$ K, but rather cools from a maximum temperature of $\sim 10^4 - 10^5$ K. We examine the possible values for $n_e$ and $f$ in these two models by calculating $f$ as a function of $n_e$ for three temperatures: $10^4$ K and $10^5$ K for cold flow accretion and $10^6$ K for cooling from the hot halo. |
In Fig. 6.. | In Fig.\ref{density}, , |
we ot f versus n», for these three temperatucs for the rauge of dpafidt at 2.=0 predicted by our open box model: LlL2«107 M. MpeP +. | we plot $f$ versus $n_e$ for these three temperatues for the range of $d\mrext/dt$ at $z=0$ predicted by our open box model: $0.4 - 1.2 \times 10^7$ $_\sun$ $\permpccu$ $^{-1}$ . |
The y-axis ou the vielt side shows logR/ρε) Corresponding to f£. where R is the radius of the relevant volume associated with au L. galaxy. | The y-axis on the right side shows $\log(R/R_{virial})$ corresponding to $f$ , where R is the radius of the relevant volume associated with an $L_*$ galaxy. |
The vertical dotted line indicates the value of n, at which the cooling time is equal to the age of the universe. | The vertical dotted line indicates the value of $n_e$ at which the cooling time is equal to the age of the universe. |
We use the approximate form for ACP) for mild enriched eas (Z= 0.1) from | We use the approximate form for $\Lambda(T)$ for mildly enriched gas $Z = 0.1$ ) from \cite{MB2004} : |
The observed accelerated: expansion of the universe (see c.8.(2))) is usually ascribed to the existence of a cosmic Iluid with a negative pressure comparable to its small but cosmologically significant energy density. | The observed accelerated expansion of the universe (see \citep{astier:2006}) ) is usually ascribed to the existence of a cosmic fluid with a negative pressure comparable to its small but cosmologically significant energy density. |
While numerous possible origins have been proposed. for this clark energy (sce (?) for a review) most of them are variations of the cosmological constant or the scalar Ποίά scenario. (7).. | While numerous possible origins have been proposed for this dark energy (see \citep{Copeland:2006wr} for a review) most of them are variations of the cosmological constant or the scalar field scenario \citep{Wetterich:1987fm}. |
These include several cifferent Kinds of couplings to matter (barvons. neutrinos. cold dark mattor(?????))). couplings to eravity (the curvature scalar. Lovelock invariants. etc (2?))) and also non-canonical kinetic terms (phantoms. tachvons. Ix-essence. ete. (??))). | These include several different kinds of couplings to matter (baryons, neutrinos, cold dark \citep{Amendola:1999er,brook,mota,mota2,Koivisto:2005nr}) ), couplings to gravity (the curvature scalar, Lovelock invariants, etc \citep{Koivisto:2006xf,Koivisto:2006ai}) ), and also non-canonical kinetic terms (phantoms, tachyons, K-essence, etc \citep{Gibbons:2002md,ArmendarizPicon:2000ah}) ). |
With such a zoo of models. 1t is important to investigate. ancl search. for. some particular feature or property of dark energy. which could be used to rule out some of (or at least distinguish among) all these possibilities. | With such a zoo of models, it is important to investigate, and search for, some particular feature or property of dark energy, which could be used to rule out some of (or at least distinguish among) all these possibilities. |
lt ds already well known in the literature that to discriminate between different classes of mioclels. it is not sullicient to consider just the background expansion. | It is already well known in the literature that to discriminate between different classes of models, it is not sufficient to consider just the background expansion. |
One has also to study the evolution of cosmological perturbations. | One has also to study the evolution of cosmological perturbations. |
General Relativity dictates that any component with its equation of state w=1. where w=| corresponds to the cosmological constant. must Iluetuate. | General Relativity dictates that any component with its equation of state $w \neq -1$, where $w=-1$ corresponds to the cosmological constant, must fluctuate. |
Pherefore a generic dark energy. component has perturbations which couple to matter perturbations. | Therefore a generic dark energy component has perturbations which couple to matter perturbations. |
These however can be small. since if one has dwzml. the component can be nearly smooth. and if the Jeans length of dark energy is large. its perturbations may be confined to very large scales only. | These however can be small, since if one has $w \approx -1$, the component can be nearly smooth, and if the Jeans length of dark energy is large, its perturbations may be confined to very large scales only. |
“Phis is typical for the usual minimally coupled. quintessence models. since their sound. speed. of perturbations. c7,,,. is equal to the light speed. c7,E=1. which sets a Large Jeans length. (22).. | This is typical for the usual minimally coupled quintessence models, since their sound speed of perturbations, $\clam$, is equal to the light speed, $\clam=1$, which sets a large Jeans length \citep{Bean:2003fb,xia}. |
Llowever. for some other dark energy candidates that is not the case (???).. and such feature might help to dilferentiate between these classes of models. | However, for some other dark energy candidates that is not the case \citep{Mota:2004pa,Bagla:2002yn,Bento:2002ps}, and such feature might help to differentiate between these classes of models. |
In addition to the w and ¢7,,,,. there is an important characteristic of a general. cosmic Eluid whieh is. its anisotropic stress σ (?7).. | In addition to the $w$ and $\clam$, there is an important characteristic of a general cosmic fluid which is its anisotropic stress $\sigma$ \citep{Hu:1998kj}. |
This vanishes for a minimally coupled. scalar field and. perfect. Haud. candidates. but is a eeneric property of realistic Huids with finite shear viscous coellicients (222)... Basically. while dw and ¢7,,,, determine respectively the background. and. perturbative pressure of the Uuicl that is rotationally invariant. σ quantifies how much the pressure of the fluid varies with direction. | This vanishes for a minimally coupled scalar field and perfect fluid candidates, but is a generic property of realistic fluids with finite shear viscous coefficients \citep{Schimd:2006pa,Brevik:2004sd,Nojiri:2005sr}.. Basically, while $w$ and $\clam$ determine respectively the background and perturbative pressure of the fluid that is rotationally invariant, $\sigma$ quantifies how much the pressure of the fluid varies with direction. |
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