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Ati26.5«107 ves. the GE goes through the second collapse when the central temperature exceeds about 2300 Ix. “Phe embryo radius drops rapidly. while the density. aud the temperature increase strongly. | At $t\approx 6.5\times 10^3$ yrs, the GE goes through the second collapse when the central temperature exceeds about 2300 K. The embryo radius drops rapidly, while the density and the temperature increase strongly. |
Phe first embryo becomes the second in our terminology. | The first embryo becomes the second in our terminology. |
Our radiation. hvdrodynamies code is not well. suited (Navakshin2010b) to follow the long cooling and contraction of the second embryo. | Our radiation hydrodynamics code is not well suited \citep{Nayakshin10b} to follow the long cooling and contraction of the second embryo. |
However. there is a body of work on the hyelrostatic contraction of very low-mass stars and Jupiter-mass planets (e.g...Crossman&Craboske1973:Graboskeetal. 1975).. which allows us to describe the process with a good degree of confidence. | However, there is a body of work on the hydrostatic contraction of very low-mass stars and Jupiter-mass planets \citep[e.g.,][]{GrossmanGraboske73,GraboskeEtal75}, which allows us to describe the process with a good degree of confidence. |
Let /» be the time of the second collapse of the embryo. | Let $t_2$ be the time of the second collapse of the embryo. |
As is well known. the embryo spends the initial contraction stages on the Llavashi track. | As is well known, the embryo spends the initial contraction stages on the Hayashi track. |
During this phase. the outward energy. transfer is dominated: by convection. | During this phase, the outward energy transfer is dominated by convection. |
The ellective temperature. Li is almost constant at logi~3.13.3 see Figs 1 in | The effective temperature, $T_{\rm eff}$ is almost constant at $\log T_{\rm eff} \sim
3.1-3.3$ [see Figs 1 in |
and the plasina velocities is determined by the following relation. Iu the above method. we can decide the electric current and the plasma velocity by the electric field aud the magnetic field in anv refereuce Ssvsteni. | and the plasma velocities is determined by the following relation, In the above method, we can decide the electric current and the plasma velocity by the electric field and the magnetic field in any reference system. |
The previous Olunus law is always confused when the exterual forces iucluded. | The previous Ohm's law is always confused when the external forces included. |
Iu this part. we will discuss the case when external forces are included in the fully ionized plasma. | In this part, we will discuss the case when external forces are included in the fully ionized plasma. |
We consider the pressure eradient. eravity or friction as the external forces f; aud. f, for the fluid of tou aud electron respectively. f;=F;vh;| {.=€Fonni.Pe | We consider the pressure gradient, gravity or friction as the external forces $f_{i}$ and $f_{e}$ for the fluid of ion and electron respectively. $
\bf{f}_{i}=\frac{\bf{F}_{i}-\bigtriangledown P_{i}}{nm_{e}}
$, $
\bf{f}_{e}=\frac{\bf{F}_{e}-\bigtriangledown P_{e}}{nm_{e}}
$. |
Tn the nearly steady case. the nn,balance equations for the ion fluid aud electron fiuid iu x. v plane will be: Iu the direction of magnetic field will be: Tere. we assume the magnetic field iu the direction of :. the electric field perpendicular to maenetic field in the direction ον aud E .Bisin the direction of y. | In the nearly steady case, the balance equations for the ion fluid and electron fluid in x, y plane will be: In the direction of magnetic field will be: Here, we assume the magnetic field in the direction of $z$, the electric field perpendicular to magnetic field in the direction $x$, and $-\bf{E}\times\bf{B}$ is in the direction of $y$. |
And equatious (32). €) 01 and (35) cau be written as: Then. from equations (37). (38). (39) aud (10) we can ect the velocities of ious aud electrons iu the plane perpendiculay to maenetic field as. Where the velocity e,=|=D?B|. | And equations (32), (33), (34) and (35) can be written as: Then, from equations (37), (38), (39) and (40) we can get the velocities of ions and electrons in the plane perpendicular to magnetic field as, Where the velocity $v_{s}=|\frac{\bf{E}\times \bf{B}}{\bf{B}^2}|$. |
So we can get. Aud if we use equation (11) subtract (12) directly we can get the velocity difference of electrons aud ious iu the direction :: By imultiphung ve. we cau ect the electric current iu the plane perpeudicular to the magnetic field. | So we can get, And if we use equation (41) subtract (42) directly we can get the velocity difference of electrons and ions in the direction $z$: By multiplying $ne$, we can get the electric current in the plane perpendicular to the magnetic field. |
reach a better agreement with the observations. | reach a better agreement with the observations. |
The rotation velocity had to be increased by up to 4 iin the innermost parts. and to be decreased by 2 iin the outer parts of the low-resolution cube. | The rotation velocity had to be increased by up to 4 in the innermost parts, and to be decreased by 2 in the outer parts of the low-resolution cube. |
The racial velocity was increased by up to 4 lin the innermost and outermost parts. | The radial velocity was increased by up to 4 in the innermost and outermost parts. |
Lhe comparison between our best model high-resolution data cube ancl the observed one is shown in Fig. 13.. | The comparison between our best model high-resolution data cube and the observed one is shown in Fig. \ref{cubes_hires}. |
With our choice of input parameters we are able to reproduce all the major features present in the observations. | With our choice of input parameters we are able to reproduce all the major features present in the observations. |
We reiterate that without the introduction of non-circular motions the good. agreement would not have been reached. due to the impossibility of reproducing simultaneously the total intensity map (Figs. | We reiterate that without the introduction of non-circular motions the good agreement would not have been reached, due to the impossibility of reproducing simultaneously the total intensity map (Figs. |
and 2)) and the central channels (Fig. 10)). | \ref{mom0_only} and \ref{mom0_opt_only}) ) and the central channels (Fig. \ref{cubes_hires_vr0}) ). |
The excellent agreement between the two cata cubes is also visible in Vig. 14. | The excellent agreement between the two data cubes is also visible in Fig. \ref{pv.3slices}, |
in which position-velocity diagrams along the major axis are compared. | in which position-velocity diagrams along the major axis are compared. |
The only feature that is not well reproduced is the northern part of the emission in channels 237.5 tto 262kms ((receding side. see Fie. 13)) | The only feature that is not well reproduced is the northern part of the emission in channels 237.5 to 262.2 (receding side, see Fig. \ref{cubes_hires}) ). |
However. since 19 corresponding emission is not present in the approaching channels. it means that it is an asvmmetric feature of the galaxy. thus it is not reproducible with our axisvmmetric moclels. | However, since the corresponding emission is not present in the approaching channels, it means that it is an asymmetric feature of the galaxy, thus it is not reproducible with our axisymmetric models. |
Bearing in mind the caveats discussed above on the use of the x7 statistics. the fact that our preferred. mode is à better reproduction of the observations (based. on the masked cube. see Section 2)) can also be seen in the smaller ovalue (47=245107 for about 103. degrees of freeclom) compared to the model data cube built from the first attempt to derive the rotation curve (Section 3. (>=3.62. 104) or compared to the model with the correct orientation angles but. no non-circular motions (Fig. 10.. | Bearing in mind the caveats discussed above on the use of the $\chi^2$ statistics, the fact that our preferred model is a better reproduction of the observations (based on the masked cube, see Section \ref{obsred}) ) can also be seen in the smaller $\chi^2$ value $\chi^2=2.45\times 10^4$ for about $\times 10^4$ degrees of freedom) compared to the model data cube built from the first attempt to derive the rotation curve (Section \ref{velfi_first}, $\chi^2=3.62\times 10^4$ ) or compared to the model with the correct orientation angles but no non-circular motions (Fig. \ref{cubes_hires_vr0}, |
y=8.97 105. | $\chi^2=3.97\times 10^4$ ). |
A similar procedure was performed. to fit the resolution cube (see Section 3.5)): the results from the low-resolution cube were used for radii larger than180".. because of its better signal-to-noise ratio in the outer parts. | A similar procedure was performed to fit the low-resolution cube (see Section \ref{model_lores}) ); the results from the low-resolution cube were used for radii larger than, because of its better signal-to-noise ratio in the outer parts. |
One of the main goals of this paper is to investigate the distribution of dark matter at large radii: we smoothed the high-resolution cube to obtain a beam of iin order to highlight the faint cülf'use emission. | One of the main goals of this paper is to investigate the distribution of dark matter at large radii; we smoothed the high-resolution cube to obtain a beam of $\times$ in order to highlight the faint diffuse emission. |
Similarly to the high resolution cube. first a. tilted- fit on the velocity. field was performed. | Similarly to the high resolution cube, first a tilted-ring fit on the velocity field was performed. |
Again. he kinematical position angle is different from. the morphological one. and it needs to be modified in order o reproduce the data cube. | Again, the kinematical position angle is different from the morphological one, and it needs to be modified in order to reproduce the data cube. |
The same procedure as above was applied. to look for the best. &eometrical and ohvsical parameters to. describe the observed. data cube. | The same procedure as above was applied to look for the best geometrical and physical parameters to describe the observed data cube. |
The parameters for the inner parts were fixed at the values erived from the high-resolution cube: based on the signal-o-noise ratio of the high-resolution cube. it was decided to nlv use the high-resolution values up to180". | The parameters for the inner parts were fixed at the values derived from the high-resolution cube; based on the signal-to-noise ratio of the high-resolution cube, it was decided to only use the high-resolution values up to. |
.. A harmonic econiposition of the velocity feld. was performed. whose results are not shown because they are similar to Fig. 11.. | A harmonic decomposition of the velocity field was performed, whose results are not shown because they are similar to Fig. \ref{harmonic}. |
The s, term in the outer parts is close to 10Lo while 16 DoD-zero es term is again interpreted as the effect ofdisk —--—uckness: the inclination is well constrained from the total ] map (Fig. 1)). | The $s_1$ term in the outer parts is close to 10, while the non-zero $c_3$ term is again interpreted as the effect of disk thickness; the inclination is well constrained from the total HI map (Fig. \ref{mom0_lores}) ). |
The comparison between the observed. ancl modelled low-resolution total HIE maps is displaved in Fig. 7.. | The comparison between the observed and modelled low-resolution total HI maps is displayed in Fig. \ref{mom0_lores}. |
Again. 1 agreement between the maps is excellent. | Again, the agreement between the maps is excellent. |
As can be seen in Figs. | As can be seen in Figs. |
15. and 10 our choice of parameters enables us to reproduce the observations in detail. | \ref{cubes_lores} and \ref{pv.60.3slices} our choice of parameters enables us to reproduce the observations in detail. |
Without the introduction of non-cireular motions such a good agreement would not have been possible. | Without the introduction of non-circular motions such a good agreement would not have been possible. |
Phe final choice of inclination. position angle. rotation velocity and radial velocity is shown in Fig. 12.. | The final choice of inclination, position angle, rotation velocity and radial velocity is shown in Fig. \ref{parametri}. |
The parameters from the low-resolution cube were sampled every ((half the beam EWIIM). so that. the low- and. high- cubes contribute to the final rotation curve with a comparable number of points. | The parameters from the low-resolution cube were sampled every (half the beam FWHM), so that the low- and high-resolution cubes contribute to the final rotation curve with a comparable number of points. |
Thanks to the exceptional extension of the LL disk. the rotation curve extends out to | Thanks to the exceptional extension of the HI disk, the rotation curve extends out to |
he non-ACN companion. | the non-AGN companion. |
Table 3/ reports the ACN fraction in each pair sample and iu the natching uupaired field sample. | Table \ref{agn_table} reports the AGN fraction in each pair sample and in the matching unpaired field sample. |
The ACN fraction in pairs exceeds that in the matching field sample or cach of the subsamples. | The AGN fraction in pairs exceeds that in the matching field sample for each of the subsamples. |
The increase in ACN raction for both the full aud volunuc-IHuited major ür Saluples is a factor of —2. a siguificauce of ~260. | The increase in AGN fraction for both the full and volume-limited major pair samples is a factor of $\sim2$, a significance of $\sim2\sigma$. |
The hieher ACN fractions in the volume-nuited samples (pair and field) compared to the Tull samples is consistent with the exchision of ower Iuninositv galaxies from the volumc-Inited salple. which are less likely to coutain an ACN, | The higher AGN fractions in the volume-limited samples (pair and field) compared to the full samples is consistent with the exclusion of lower luminosity galaxies from the volume-limited sample, which are less likely to contain an AGN. |
It is useful to segregate galaxies iuto two categories: carly-type galaxies. which are ecucrally eas-poor and have little or no active star formation. and late-tvpe. which coutain more gas and have vouug stellar populations. | It is useful to segregate galaxies into two categories: early-type galaxies, which are generally gas-poor and have little or no active star formation, and late-type, which contain more gas and have young stellar populations. |
The eas-vich svstcus are potenti:]v susceptible to tidally triggered star formation im major interactions (c.g.Milos2006:DiMatteoetal. 2007). | The gas-rich systems are potentially susceptible to tidally triggered star formation in major interactions \citep[e.g.][]{mihos+hern96,tissera,cox06,dimatteo07}. |
. Tuteractious between eas-poor galaxies ("dry mergers) produce little or no star formation activity. although. thev contribute substantially to the build-ap of massive ealaxies (e.g.Tranetal.2005:Vau Dokkuu2005:Cattaneoetal. 2008). | Interactions between gas-poor galaxies (“dry mergers”) produce little or no star formation activity, although they contribute substantially to the build-up of massive galaxies \citep[e.g.][]{tran05,vandokkum05,cattaneo08}. |
.. Our analysis of star formation activity iu pair galaxies focuses on the late-tvpe svstenis. | Our analysis of star formation activity in pair galaxies focuses on the late-type systems. |
There ire à uuuber of classification schemes to separate the carly aud late-type galaxy populations. | There are a number of classification schemes to separate the early and late-type galaxy populations. |
Photometric cdiscrininuauts imclude color. concentration. and absolute maeuitucde (e.g.Stratevaotal.2001:Ikauffiuiaunetal. 20035)... aud spectroscopic nethods include D,1000. and 15 absorptio- (Ikauffinaàunetal.2003a). | Photometric discriminants include color, concentration, and absolute magnitude \citep[e.g.][]{strateva01,kauff03_54}, and spectroscopic methods include $D_n4000$ and $\delta$ absorption \citep{kauff03_33}. |
. The D,,1000 indicator discriminates by stella population age. | The $D_n4000$ indicator discriminates by stellar population age. |
At waveleugtlis Auer than £000.A.. metal lines iu low mass stars absorb the light aud cause a “break” in he προςπα, | At wavelengths bluer than 4000, metal lines in low mass stars absorb the light and cause a “break” in the spectrum. |
As the stellar population ages aud he massive. hot stars die off. D,,1000. increases unonotouicallv with time. | As the stellar population ages and the massive, hot stars die off, $D_n4000$ increases monotonically with time. |
Ikauffinaunetal.(20034) use stellar populatiμαi models to show that ealaxies with D,LOOSs15 lave voung stellar oopulatious (=1 Cr). | \citet{kauff03_33} use stellar population models to show that galaxies with $D_n4000 \lesssim 1.5$ have young stellar populations $\lesssim 1$ Gyr). |
Motallicity has a strong effect on the the value of D,,LOOO ouly after 1 Cir ost a burst of star formation (seeFie.2iuIautt- 2003a).. | Metallicity has a strong effect on the the value of $D_n4000$ only after 1 Gyr past a burst of star formation \citep[see Fig. 2 in][]{kauff03_33}. |
Measurement of D,1000. is insensitive to ealaxv reddening. | Measurement of $D_n4000$ is insensitive to galaxy reddening. |
Vorganietal.(2008) use D,,1000 to separate spectroscopic earlv-tvpe galaxies from late-tvpoe ealaxies at a dividing line of D,,1000=1.5 in their analvsis of galaxv stellar mass asseniblv iu the VIMOS VLT Deep Survey. | \citet{vergani08} use $D_n4000$ to separate spectroscopic early-type galaxies from late-type galaxies at a dividing line of $D_n4000 = 1.5$ in their analysis of galaxy stellar mass assembly in the VIMOS VLT Deep Survey. |
Mignolietal.(2005) also use ID1000=1.6 as a dividiug line (different definition of {21000 from Bruzual1983: ratio offlux in bands LO50-1250 tto. 3150-3950. Aj) along with other spectral nieasurenients fo classify ealaxics in the K20 survey, a near-IR selected redshift survey. | \citet{mignoli05} also use $D4000=1.6$ as a dividing line (different definition of $D4000$ from \citealp{bruzual83}: ratio offlux in bands 4050-4250 to 3750-3950 ) along with other spectral measurements to classify galaxies in the K20 survey, a near-IR selected redshift survey. |
We follow this approach. | We follow this approach. |
The distribution of D,,000 is bimodal. separating ealaxv populations comiuated by old stars from systems with recent star formation. | The distribution of $D_n4000$ is bimodal, separating galaxy populations dominated by old stars from systems with recent star formation. |
Figure 9 shows the distribution of D,1000 for the 6.611 ealaxies (both pair aud unupaired) at 2=0.0800.376 with a robust D,1000. measurement. | Figure \ref{hist_d4000} shows the distribution of $D_n4000$ for the 6,644 galaxies (both pair and unpaired) at $z=0.080-0.376$ with a robust $D_n4000$ measurement. |
The peaks are at —1.15 and ~1.75. | The peaks are at $\sim 1.15$ and $\sim 1.75$. |
Ikaufftinaunetal. similarly observe a bimodal distribution in D,LOOO. with peaks at 1.30 and 1.85 in their sample of ~100.000 ealaxics in the SDSS. | \citet{kauff03_33} similarly observe a bimodal distribution in $D_n4000$ , with peaks at 1.30 and 1.85 in their sample of $\sim100,000$ galaxies in the SDSS. |
We choose the imuunmuimn between the ποσα distribution as our dividing line. D,,1000=1.11 (Figure 9)). | We choose the minimum between the bimodal distribution as our dividing line, $D_n4000 = 1.44$ (Figure \ref{hist_d4000}) ). |
We refer to galaxies with D,,LOOΊντι as "low D,1000 ealaxics and those with D,1000>L.1E a8 whieh” D,000. galaxies. | We refer to galaxies with $D_n4000 \leq 1.44$ as “low” $D_n4000$ galaxies and those with $D_n4000>1.44$ as “high” $D_n4000$ galaxies. |
Our analysis of star formation activity includes the low D,,LOOO ealaxy in a mixed pair. but does not require both galaxies to have D,,1000.<1.LL. | Our analysis of star formation activity includes the low $D_n4000$ galaxy in a mixed pair, but does not require both galaxies to have $D_n4000 < 1.44$. |
The emission line fraction is a stroug fuuctio- of D,1000. | The emission line fraction is a strong function of $D_n4000$. |
Fieure 10 shows the steep declan iu the fraction of euission line galaxies betweeu 121000=1.5.1.5. | Figure \ref{emfracd4} shows the steep decline in the fraction of emission line galaxies between $D_n4000 = 1.3-1.5$. |
The enisson line galaxieshave EW(Ila)=3 or EW([O IIl)»3. | The emission line galaxieshave $\alpha) \geq 3$ or EW([O $\geq 3$. |
Thus seeregatiugby D,,000 Is reasonable aud comespouds well to segregatiug by the preseuce of emisso- lines. | Thus segregating by $D_n4000$ is reasonable and corresponds well to segregating by the presence of emission lines. |
We also compare the use of D,1000 for ealaxyv classification with color. another widely used indicator of galaxy type. | We also compare the use of $D_n4000$ for galaxy classification with color, another widely used indicator of galaxy type. |
Figure Ll shows the bimodal distribution of rest-frame SDSS (yg r) color versus D, 1000. | Figure \ref{clr_d4000} shows the bimodal distribution of rest-frame SDSS $g-r$ ) color versus $D_n4000$ . |
We compute the rest-frame SDSS (g r) color using thekle correction determined by Annis(2001) from the Pegase code (LeBorene&Rocca-Voliunerage 2002)... | We compute the rest-frame SDSS $g-r$ ) color using thek+e correction determined by \citet{annis01} from the Pegase code \citep{pegase}. . |
Th separation by D,,1000 is more sharply defined than that of rest-frame color. | The separation by $D_n4000$ is more sharply defined than that of rest-frame color. |
The rest-frame ealaxy colors are affecte by reddening. aud depend o- the noisy k|e corrections. | The rest-frame galaxy colors are affected by reddening, and depend on the noisy k+e corrections. |
D,,I000 provides a | $D_n4000$ provides a |
4-meter in 0.35" seeing. | 4-meter in $^{\prime\prime}$ seeing. |
Figure 10 compares the weak lensing sensitivity of the two fields. | Figure \ref{fig:sensitive2.ps} compares the weak lensing sensitivity of the two fields. |
The curves show the rest frame line-of sieht velocity dispersion lor a svstem which could be detected al av =3.7 a4 each vredshift. | The curves show the rest frame line-of sight velocity dispersion for a system which could be detected at a $\nu = 3.7$ at each redshift. |
AC every redshift. the Φρα observations have the potential to vield a detection al a lower line-ol-sieht velocity dispersion. | At every redshift, the Subaru observations have the potential to yield a detection at a lower line-of-sight velocity dispersion. |
The maximum sensitivitv is shifted toward greater redshift for the Subaru data. | The maximum sensitivity is shifted toward greater redshift for the Subaru data. |
Table 2. shows the substantial difference in the number of resolved sources for the (wo weak lensing maps: 22.4 ? for the DLS and 35 ? [for the GTO2deg? field. | Table \ref{tbl:Maps} shows the substantial difference in the number of resolved sources for the two weak lensing maps: 22.4 $^{-2}$ for the DLS and 35 $^{-2}$ for the $^2$ field. |
Table 2. lists other relevant parameters for the (wo survevs. | Table \ref{tbl:Maps} lists other relevant parameters for the two surveys. |
In constructing Table 2. we have reevaluated ihe DLS F2 results of Geller et al. ( | In constructing Table \ref{tbl:Maps} we have reevaluated the DLS F2 results of Geller et al. ( |
2010) using a v=3.7 rather than av =3.5 threshold lor selecting significant peaks in the weak lensing map. | 2010) using a $\nu = 3.7$ rather than a $\nu =3.5$ threshold for selecting significant peaks in the weak lensing map. |
We have also included (he Mivazaki et al. ( | We have also included the Miyazaki et al. ( |
2007) procedure for removing unreliable. but apparently. high significance weak lensing peaks. | 2007) procedure for removing unreliable, but apparently high significance weak lensing peaks. |
Figure 16. also shows (he rest frame line-of sight velocity dispersions for svstems in the [oregrouncd redshift surveys corresponding to weak lensing peaks with v> 3.7 in both survevs. | Figure \ref{fig:sensitive2.ps} also shows the rest frame line-of sight velocity dispersions for systems in the foreground redshift surveys corresponding to weak lensing peaks with $\nu > $ 3.7 in both surveys. |
The most telling ancl suggestive difference between the adimittedly small samples of clusters in the two surveys is that the demonstration that. as expected [rom (he sensitivity curves. the Subaru map cleanly detects clusters with rest [rame line-ol sieht. velocity significantly below the limits of the DLS sensitivity al (he same redshilt. | The most telling and suggestive difference between the admittedly small samples of clusters in the two surveys is that the demonstration that, as expected from the sensitivity curves, the Subaru map cleanly detects clusters with rest frame line-of sight velocity significantly below the limits of the DLS sensitivity at the same redshift. |
For example. the Subaru lensing map cleanly detects a cluster at z=0.540 with a rest Game line-ol-sight velocity dispersion of ~600 kms |: at a similar redshift the DLS detects a svstem with a much larger velocity dispersion of ~900 km ! at av~3.7. | For example, the Subaru lensing map cleanly detects a cluster at $ z = 0.540$ with a rest frame line-of-sight velocity dispersion of $\sim 600$ km $^{-1}$; at a similar redshift the DLS detects a system with a much larger velocity dispersion of $\sim 900$ km $^{-1}$ at a $\nu \sim 3.7$. |
This comparison indicates that. in principle. observations with good seeing can be the basis for a weak lensing measurement of (he cluster mass function for masses >1.7x10! AL. (rest [rame line-ol-sight velocity dispersion >600 km ! ) at redshifts 0.2Xz<0.6. | This comparison indicates that, in principle, observations with good seeing can be the basis for a weak lensing measurement of the cluster mass function for masses $\gtrsim 1.7\times 10^{14}$ $_\odot$ (rest frame line-of-sight velocity dispersion $\gtrsim 600$ km $^{-1}$ ) at redshifts $0.2 \lesssim z \lesssim 0.6$. |
Understanding both the completeness and the ellicieney of weak lensing map cluster detections are. of course. crucial to their use as cosmological tools for measuring e.g. the cluster mass function. | Understanding both the completeness and the efficiency of weak lensing map cluster detections are, of course, crucial to their use as cosmological tools for measuring e.g. the cluster mass function. |
The DLS weak lensing cluster sample is ~50% complete relative to a ssmple drawn from the foreground. SIIELS redshilt survey. (Geller οἱ al. | The DLS weak lensing cluster sample is $\sim 50$ complete relative to a sample drawn from the foreground SHELS redshift survey (Geller et al. |
2010). | 2010). |
The GTO2dee? cluster catalog is complete (Section 4.2)). but the sample is so small that we cannot draw a general conclusion. | The $^2$ cluster catalog is complete (Section \ref{clusters}) ), but the sample is so small that we cannot draw a general conclusion. |
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