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1996) demonstrated that the detection of substructure was more efficient when MN, "Sepends on the number of cluster members. | 1996) demonstrated that the detection of substructure was more efficient when $N_{local}$ depends on the number of cluster members. |
They proposed NyNa. Νιμ being the number of cluster members. | They proposed $N_{local}=\sqrt{N_{gal}}$, $N_{gal}$ being the number of cluster members. |
Numerical Vsimulations also show that the DS test effectively detects substructures in clusters with more than 30 galaxies. | Numerical simulations also show that the DS test effectively detects substructures in clusters with more than 30 galaxies. |
Here. we have studied the substructure of those galaxy clusters with Nig)>30 by using Nj=νΝιμ. | Here, we have studied the substructure of those galaxy clusters with $N_{gal}> 30$ by using $N_{local}=\sqrt{N_{gal}}$. |
Biviano et al. ( | Biviano et al. ( |
2002) computed 0; for each galaxy as the average of the 0-values of its Niegi-| neighbours. | 2002) computed $\delta_{i}$ for each galaxy as the average of the $\delta$ -values of its $_{local}$ -1 neighbours. |
We did not find any important differences when carrying out this smoothing. so we finally decided to work with the unsmoothed 0-values of each galaxy. | We did not find any important differences when carrying out this smoothing, so we finally decided to work with the unsmoothed $\delta$ -values of each galaxy. |
In order to test the code used for detecting substructure in our galaxy clusters we ran several Monte Carlo simulations of clusters. | In order to test the code used for detecting substructure in our galaxy clusters we ran several Monte Carlo simulations of clusters. |
These simulated clusters did not hold substructure and were used in order to test the false positives cases detected by our code. Le. clusters with substructure reported by the test but without substructure. | These simulated clusters did not hold substructure and were used in order to test the false positives cases detected by our code, i.e. clusters with substructure reported by the test but without substructure. |
We considered all our clusters with more than 30 galaxies in ro<2raq9. and substituted their velocity distribution by a Gaussian distribution with the same v, and σι as the real clusters. | We considered all our clusters with more than 30 galaxies in $r<2 r_{200}$, and substituted their velocity distribution by a Gaussian distribution with the same $_{c}$ and $\sigma_{c}$ as the real clusters. |
This kind of simulation breaks the possible correlation between galaxy position and velocity. i.e. possible substructure. | This kind of simulation breaks the possible correlation between galaxy position and velocity, i.e. possible substructure. |
We ran the DS test on these simulated clusters as for the real ones. giving no signal of substructure in all cases. | We ran the DS test on these simulated clusters as for the real ones, giving no signal of substructure in all cases. |
This means that our test does not detect cases of false positives in the study of the global substructure in clusters. | This means that our test does not detect cases of false positives in the study of the global substructure in clusters. |
We have also ran another set of MC simulations of clusters in order to analyse the efficiency of the DS test for detecting substructure. | We have also ran another set of MC simulations of clusters in order to analyse the efficiency of the DS test for detecting substructure. |
In contrast to the previous set of simulations. these simulated clusters held substructure. | In contrast to the previous set of simulations, these simulated clusters held substructure. |
Thus. the new simulated clusters had two populations of galaxies: one with no correlation between position and velocities showing a Gaussian velocity distribution (we call this a relaxed population) and aother of galaxies showing substructure (we call this an unrelaxed population). | Thus, the new simulated clusters had two populations of galaxies: one with no correlation between position and velocities showing a Gaussian velocity distribution (we call this a relaxed population) and another of galaxies showing substructure (we call this an unrelaxed population). |
The surface distribution of the galaxies | The surface distribution of the galaxies |
astronomical comunity. | astronomical community. |
An analysis of these spectra with details on the reduction procedure and a log of observations can be found in the original paper. | An analysis of these spectra with details on the reduction procedure and a log of observations can be found in the original paper. |
This QSO is one of the brightest known also by virtue of eravitatioua lensing maenification. | This QSO is one of the brightest known also by virtue of gravitational lensing magnification. |
The image of this QSO reveals two commpoucuts of similar iuteusity separated by about 0", (Irwin et al Ellison et al describe the absorber a ται. and identify the high ionization species of SHV and CTV associated with the svstem. | The image of this QSO reveals two components of similar intensity separated by about $0''.4$ (Irwin et al \nocite{Irwin}
Ellison et al \nocite{E99} describe the absorber at $z_{\rm abs}$ =3.514 and identify the high ionization species of SiIV and CIV associated with the system. |
A total of five components are resolved in SiTV and two in CTV. with both ious showing a main component at redshift Zaha=3.5138, | A total of five components are resolved in SiIV and two in CIV, with both ions showing a main component at redshift $_{\rm abs}=3.5138$. |
Our attention was caught bv the feature on the blue wing of the Lya profile of the displaved in Fie. | Our attention was caught by the feature on the blue wing of the $\alpha$ profile of the displayed in Fig. |
lof Ellison et al.. which we suspected to be a D line. | 4 of Ellison et al \nocite{E99}, which we suspected to be a D line. |
In acelition we have identified features due to CIT 133.15A.. Sill 1260. aand Silll 1206.5 (see Fie. | In addition we have identified features due to CII 1334.5, SiII 1260.4 and SiIII 1206.5 (see Fig. |
1). | 1). |
The CIT aud Sill lines are very faint am the detection has been possible thauks to the quality of the TIRES spectra at this waveleneth (S/N ~ 125). | The CII and SiII lines are very faint and the detection has been possible thanks to the quality of the HIRES spectrum at this wavelength (S/N $\sim$ 125). |
The SUT 1193.3 aabsorption mav also be preseut. albeit strongly bleudec with a lower redshift Lye absorber. | The SiII 1193.3 absorption may also be present, albeit strongly blended with a lower redshift $\alpha$ absorber. |
Ly. falls iu the available spectral rauge but is. uufortunatelv. severely blended with local Lya clouds (Fig. | $\beta$ falls in the available spectral range but is, unfortunately, severely blended with local $\alpha$ clouds (Fig. |
1. panel ο). | 1, panel $c$ ). |
The CID profile shows a double compoueut with the main absorption occuring at redshift z4,,223.5127I. namely the redshift at which the main features of CIX (Zany= 351380) ancl SUV (zai.= 3.51376) also occur. as measted by EEllison et al (1999). | The CII profile shows a double component with the main absorption occurring at redshift $_{\rm abs}$ =3.51374, namely the redshift at which the main features of CIV $_{\rm abs}=3.51380$ ) and SiIV $_{\rm abs}=3.51376$ ) also occur, as measured by \nocite{E99}E Ellison et al (1999). |
A third CTI 1331 absorption at about |200 kins 1 Is assoclated to another Ίνα cloud redwards of our svstem. | A third CII 1334 absorption at about +200 km $^{-1}$ is associated to another $\alpha$ cloud redwards of our system. |
The CID and Sill ious are better tracers of neutral hvdroseu (aud deuteruni) than higher ionization species iu Lxiuan lait svstenis (PProchaska. 1999). | The CII and SiII ions are better tracers of neutral hydrogen (and deuterium) than higher ionization species in Lyman limit systems \nocite{proch}P Prochaska, 1999). |
To fit the observed profiles we used the package (Foutana Ballester 1995) inMIDAS. | To fit the observed profiles we used the package (Fontana Ballester 1995) \nocite{fontana}
in. |
The results for the analysis are reported in Table 1 aud the fits are shown in Fie. | The results for the analysis are reported in Table 1 and the fits are shown in Fig. |
l. | 1. |
We model the Lya absorption profile | We model the $\alpha$ absorption profile |
By using our results for the GCLF turnover magnitudes. along with previous results for several other gE galaxies from the literature a classic "EIubble Diagram? of redshilt ez against the apparent. turnover magnitude V" of the GCLF can be constructed (see and Navelaarsetal.(2000) for the first use of this diagram for GCLFs). | By using our results for the GCLF turnover magnitudes, along with previous results for several other gE galaxies from the literature a classic “Hubble Diagram” of redshift $cz$ against the apparent turnover magnitude $V^0$ of the GCLF can be constructed (see \citet{Har01} and \citet{Kav00} for the first use of this diagram for GCLFs). |
IIubble's law states ez=Hyd. | Hubble's law states $cz=H_0 d$. |
This can be rewritten as: where My. is the GCLF turnover huninositygalaxies. and If, is expressed in the usual units of kms.+Mpe.+. | This can be rewritten as: where $M_V^0$ is the GCLF turnover luminosity, and $H_0$ is expressed in the usual units of ${\rm km \; s^{-1} \;
Mpc^{-1}}$. |
Plotting log(ez) against the apparent magnitude V" gives a straight line of slope 0.2. and a zeropoint which contains //, and AL}. | Plotting $cz$ ) against the apparent magnitude $V^0$ gives a straight line of slope 0.2, and a zeropoint which contains $H_0$ and $M_V^0$. |
The available data for a total of eleven BCG galaxies or groups are listed in Table and plotted in Figure 7.. | The available data for a total of eleven BCG galaxies or groups are listed in Table \ref{TO_table} and plotted in Figure \ref{hub_plot}. |
The values for the Virgo and Fornax clusters are (he weighted mean (V5 values of the individual galaxies(see Kavelaarsetal. (2000))). | The values for the Virgo and Fornax clusters are the weighted mean $\left<V^0\right>$ values of the individual galaxies(see \citet{Kav00}) ). |
The mean radial velocities of each galaxy or group are taken [from Faberetal.(1989).. Qirardiοἱal. (1993). ]Óluchra (1938).. DinggeliPopeseu&Tanunann (1993)... Παινοἱal. (1996)... Dunn (1996).. and Laueretal.(1998). | The mean radial velocities of each galaxy or group are taken from \citet{Fab89}, \citet{Gir93}, \citet{Huc88}, \citet{BPT93}, , \citet{Ham96}, \citet{CoDu96}, and \citet{Lau98}. |
. Also. the recession velocities ez for the target galaxies asstune a Local Group infall to Virgo of 250+100kms.! (e.g.. Fordetal.(1996): (1996): Jerjen&Tanumann (1993).. among many others).The four galaxies we have investigated here fall well within the pattern established by the others (rom Virgo at low redshift out to Coma at the highest redshift). | Also, the recession velocities $cz$ for the target galaxies assume a Local Group infall to Virgo of $250 \pm 100 \; {\rm km \; s^{-1}}$ (e.g., \citet{Ford96}; \citet{Ham96}; \citet{JeTa93}, among many others).The four galaxies we have investigated here fall well within the pattern established by the others (from Virgo at low redshift out to Coma at the highest redshift). |
A value of //, near 70 km ! I. which represents aconsensus of recent determinations (Freedman 2001).. along with A/))=—7.33(Iavelaarsetal.2000:Harris2001).. matches the total range of points within their measurement uncertainties. | A value of $H_0$ near $\simeq 70$ km $^{-1}$ $^{-1}$ , which represents aconsensus of recent determinations \citep{Fre01}, , along with $M_V^0 = -7.33$\citep{Kav00,
Har01},, matches the total range of points within their measurement uncertainties. |
Spectroscopic signatures of flares on Al dwaarls have been sporadically explored. mostly al oplical wavelengths. since the pioneering observations of UV. Celi by Jov IIuniason (1949). | Spectroscopic signatures of flares on M dwarfs have been sporadically explored, mostly at optical wavelengths, since the pioneering observations of UV Ceti by Joy Humason (1949). |
Principal features observed during flares include strong continuum radiation rising toward the blue and near-ultraviolet: broad. enhanced hydrogen Dalmer series emission lines: | Principal features observed during flares include strong continuum radiation rising toward the blue and near-ultraviolet; broad, enhanced hydrogen Balmer series emission lines; |
and Ternquist 2000). | and Hernquist 2000). |
The present application to the P?—P diagram shows no indication of an ellect on the pulsar distribution due to depletion of the disk before pulsar turnoff. except possibly [or 1e Oldest pulsars. well advanced in their evolution on propeller spin-clown tracks. dropping lown to pure dipole tracks near death lines. | The present application to the $P-\dot{P}$ diagram shows no indication of an effect on the pulsar distribution due to depletion of the disk before pulsar turnoff, except possibly for the oldest pulsars, well advanced in their evolution on propeller spin-down tracks, dropping down to pure dipole tracks near death lines. |
This max indicate that there is no separate disk imescale. | This may indicate that there is no separate disk timescale. |
Disks interacting wilh pulsars via propeller torques may be driven by these torques to evolve on the propeller spin-down timescales of (he pulsars. | Disks interacting with pulsars via propeller torques may be driven by these torques to evolve on the propeller spin-down timescales of the pulsars. |
These problems. as well as je assumption that the disk remains attached to the light ¢vlinder require huther work on je coupling between the disk and the pulsar magnetosphere. | These problems, as well as the assumption that the disk remains attached to the light cylinder require further work on the coupling between the disk and the pulsar magnetosphere. |
A Monte-Carlo simulation of je P— diagram will test (he present model with independent random distributions of 2 and M and selection effects. | A Monte-Carlo simulation of the $P-\dot{P}$ diagram will test the present model with independent random distributions of $B$ and $\dot{M}$ and selection effects. |
We predict that some pulsars have braking indices <2. like the Vela pulsar (though the n=1440.2 braking index measurement for (his pulsar should be treated with some caution as its timing behaviour is dominated bv intergliteh relaxation). | We predict that some pulsars have braking indices $<2$, like the Vela pulsar (though the $n=1.4\pm 0.2$ braking index measurement for this pulsar should be treated with some caution as its timing behaviour is dominated by interglitch relaxation). |
About 2/3 of all pulsars will be on the propeller spin-down branch. with p«2. at P>Hj. | About 2/3 of all pulsars will be on the propeller spin-down branch, with $n<2$, at $P>P_0$. |
Half of these pulsars will have negative braking indices. —1<n«0. | Half of these pulsars will have negative braking indices, $-1<n<0$. |
Measurement of braking indices is unfortunately very difficult since the timing behaviour is dominated by noise (Bavkal οἱ al 1999) and it seems also by intergliteh recovery (Alpar anc Baykal 2001). | Measurement of braking indices is unfortunately very difficult since the timing behaviour is dominated by noise (Baykal et al 1999) and it seems also by interglitch recovery (Alpar and Baykal 2001). |
The model also predicts very small numbers ANxP. "ofthe oldest pulsars. which have long periods and large P values. | The model also predicts very small numbers $\Delta N \propto P^{-3}$ of the oldest pulsars, which have long periods and large $\dot{P}$ values. |
These pulsars would give magnetic fields B.—LORLoh, perhaps 10! G with the combined dipole-propeller spin-down model. | These pulsars would give magnetic fields $B_{\bot} \sim 10^{12}-10^{13}$, perhaps $10^{14}$ G with the combined dipole-propeller spin-down model. |
For these pulsars the pure dipole spin-down model would vield higher fields. extending into the magnetar range. and voung ages. | For these pulsars the pure dipole spin-down model would yield higher fields, extending into the magnetar range, and young ages. |
Thus kinematic age measurements of pulsars wil long periods and large n can disünguish between pure dipole spindown and propeller spindown. | Thus kinematic age measurements of pulsars wih long periods and large $\dot{P}$ can distinguish between pure dipole spindown and propeller spindown. |
Propeller spin-down also increases (he rate of energy. dissipation in pulsars aud in neutron stars evolving under propeller torques after pulsar activity is over. | Propeller spin-down also increases the rate of energy dissipation in pulsars and in neutron stars evolving under propeller torques after pulsar activity is over. |
We thank O. IL. Guseinov and U.. Ertan for discussions and S. C. Inna for technical support. | We thank O. H. Guseinov and Ü.. Ertan for discussions and S. Ç.. İnnam for technical support. |
We thank TUBBITTAN. the Scientific and Technical Research Council of Turkey. for support through TBAG-CCG4 and through the BDP program [or doctoral research. | We thank TÜBBİTTAK, the Scientific and Technical Research Council of Turkey, for support through TBAG-ÇGG4 and through the BDP program for doctoral research. |
AA EY thank TUDDITTAK lor graduate student scholarships. | AA EY thank TÜBBİTTAK for graduate student scholarships. |
MAA thanks the Turkish Academy of Sciences and University [or research support. | MAA thanks the Turkish Academy of Sciences and University for research support. |
the mass larger than 104A£.. the difference between CG20 results aud analytical ones are relatively large. | the mass larger than $10^{11} M_{\odot}$, the difference between CG2048 results and analytical ones are relatively large. |
Towever.18 these halos are rare objects in CG2018 run. and the fact might affect the results in some degrees. | However, these halos are rare objects in CG2048 run, and the fact might affect the results in some degrees. |
We can conclude that the shallowing slope of the mass-concentration relation naturally cmoerees from the nature of the power spectrum of iitial density fluctuations. | We can conclude that the shallowing slope of the mass-concentration relation naturally emerges from the nature of the power spectrum of initial density fluctuations. |
The slope is slieltly shallower than that of exp for lueger halos. | The slope is slightly shallower than that of $c_{\rm NFW}$ for larger halos. |
For the case of expyw. tho slope is around 0.10 for relaxed halos aud 011 for all halos2007). | For the case of $c_{\rm NFW}$ , the slope is around $-0.10$ for relaxed halos and $-0.11$ for all halos. |
. On the other haud. for the (2015 simulation. the slope is arouik 0.07 for halos with the mass LOMAS... and 0.06 for halos with the mass LO?AL... | On the other hand, for the CG2048 simulation, the slope is around $-0.07$ for halos with the mass $10^{10}M_{\odot}$, and $-0.06$ for halos with the mass $10^{9}M_{\odot}$. |
Note that one overestimates the ceutral deusity of halos if one estimates the concentration of dwart-sized halos by extrapolating the mass-concentration relation of galaxy or cluster-sized halos. | Note that one overestimates the central density of halos if one estimates the concentration of dwarf-sized halos by extrapolating the mass-concentration relation of galaxy or cluster-sized halos. |
Figure 9. shows the probability distribution functious of the concentration parameter at 2=0 in two different mass ranges. | Figure \ref{fig:cvmax_pba} shows the probability distribution functions of the concentration parameter at $z=0$ in two different mass ranges. |
Both shapes are well fitted. bv the log-normal distributions. We find logey=1.050.00.121 for halos with the mass of 5.0«LOSAL..<nassAL107AL... logey=1.022.00.128 for Lalos with the of OPAL.«M1019AJ... and logey=0.965.700.125 for halos with the mass of 1039MxOTA... | Both shapes are well fitted by the log-normal distributions, We find $\log c_0 = 1.050, \sigma = 0.124$ for halos with the mass of $5.0 \times 10^8 M_{\odot} \le M < 10^9 M_{\odot}$, $\log c_0 = 1.022, \sigma = 0.128$ for halos with the mass of $10^9 M_{\odot} \le M < 10^{10} M_{\odot}$, and $\log c_0 = 0.965, \sigma = 0.125$ for halos with the mass of $10^{10} M_{\odot} \le M < 10^{11} M_{\odot}$. |
The dimensionless spin parameter is a good parametcr to quantity the rotation of a halo. | The dimensionless spin parameter is a good parameter to quantify the rotation of a halo. |
One often uses the spin parameter defined iu(2001a).. where AV. HR. V. and Jo ds the virial mass of the halo. radius. rotational velocity at A. aud total angular monmentun inside 1. | One often uses the spin parameter defined in, where $M$, $R$, $V$, and $J$ is the virial mass of the halo, radius, rotational velocity at $R$, and total angular momentum inside $R$. |
The distribution. the dependence on the halo mass. and the evolution have been studied by a number of works2010).. | The distribution, the dependence on the halo mass, and the evolution have been studied by a number of works. |
The spin of galaxy sized halos are well studied by using the results of sufficient resolution simulations. | The spin of galaxy sized halos are well studied by using the results of sufficient resolution simulations. |
However. we do not understand those of dwarf galaxy sized halos. | However, we do not understand those of dwarf galaxy sized halos. |
The spin distribution of those halos at only high redshifts are studied bv the result of lieh resolution simulation2008). | The spin distribution of those halos at only high redshifts are studied by the result of high resolution simulation. |
Tere. we extend the spin istributions at ;=0 to dwarf galaxy sized halos (down o LOCAL.) in the same wav as the concentration. | Here, we extend the spin distributions at $z=0$ to dwarf galaxy sized halos (down to $10^8 M_{\odot}$ ) in the same way as the concentration. |
First. we deteriune the nmiuiuun number of particles oe 1a halo necessary to reliably determine the spin as doue or the concentration. | First, we determine the minimum number of particles in a halo necessary to reliably determine the spin as done for the concentration. |
Figure 109 shows the normalized ifference of average spin between the CCG2018 ruu aud he €C€6$512 run as a functiou of halo mass. | Figure \ref{fig:m-spin_res} shows the normalized difference of average spin between the CG2048 run and the CG512 run as a function of halo mass. |
We cau see hat the difference is ~0.05 for halo mass larecr than SU«109AZ... | We can see that the difference is $\sim 0.05$ for halo mass larger than $8.0 \times 10^9 M_{\odot}$. |
For halo mass less than 840«10AZ... the ifercuce is laree. | For halo mass less than $8.0 \times 10^9 M_{\odot}$, the difference is large. |
In the C512 run. a halo of mass SO\107AZ. contains 1000 particles. | In the CG512 run, a halo of mass $8.0 \times 10^9 M_{\odot}$ contains $\sim 1000$ particles. |
So we couchide hat we need 1000 particles to reliably determine the concentration. | So we conclude that we need $\sim 1000$ particles to reliably determine the concentration. |
For the (2015 run. the reliability uat is L28«10M. | For the CG2048 run, the reliability limit is $1.28 \times 10^8M_{\odot}$. |
Figure ll shows the median. aud first third quantiles of the spin parameter as a function of the virial mass of the halo. | Figure \ref{fig:m-spin} shows the median, and first third quantiles of the spin parameter as a function of the virial mass of the halo. |
Áppareutly. we can see the spin parameter is independent of the mass down to 105AZ. as pointed out for larger halos in previous works2010). | Apparently, we can see the spin parameter is independent of the mass down to $10^8 M_{\odot}$ as pointed out for larger halos in previous works. |
.. The mecdian valueis 0.0336. | The median valueis 0.0336. |
Figure 12 shows the probability distribution functions of the spin parameter at 2=0in fwo different mass ranges. | Figure \ref{fig:spin_pba} shows the probability distribution functions of the spin parameter at $z=0$in two different mass ranges. |
Both distributions are almost identical. | Both distributions are almost identical. |
The distributions are well fitted by the log-normal distributions. We find logAy= Ll7.00.308 for halos with the mass of L28&LOSAL.<AL«c109 AZ... for halos with the mass of 107AL.< (PAL... and logAy=1.152.60.277 fox halos | The distributions are well fitted by the log-normal distributions, We find $\log \lambda_0 = -1.477, \sigma = 0.308$ for halos with the mass of $1.28 \times 10^8 M_{\odot} \le M < 10^9 M_{\odot}$ , $\log \lambda_0 = -1.480, \sigma = 0.288$ for halos with the mass of $10^9 M_{\odot} \le M < 10^{10} M_{\odot}$ , and $\log \lambda_0 = -1.472, \sigma = 0.277$ for halos |
tests are independent of distance. still a very uncertain quantity for these objects (see. e.g.. Launhardt Henning 1997). | tests are independent of distance, still a very uncertain quantity for these objects (see, e.g., Launhardt Henning 1997). |
Finally. in $6) we interpret our results in the framework of an evolutionary sequence of core shapes as a function of time. | Finally, in \ref{sec-discuss} we interpret our results in the framework of an evolutionary sequence of core shapes as a function of time. |
In this picture. core morphology at early times is primarily determined by the nature of the surrounding nonisotropic mass distribution. | In this picture, core morphology at early times is primarily determined by the nature of the surrounding nonisotropic mass distribution. |
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