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In some images. clusters of galaxies are clearly visible. | In some images, clusters of galaxies are clearly visible. |
The increasing numbers of background sources at 2 23. combined with the cifliculty of accurate size measurements below this magnitude. makes efficient rejection of contaminants challenging. | The increasing numbers of background sources at $z \ga
23$ , combined with the difficulty of accurate size measurements below this magnitude, makes efficient rejection of contaminants challenging. |
For gEs this is | For gEs this is |
aud galaxies with (B/T);,, > 0.75 do uot require a bulee|disk model to fit their light. profiles. | and galaxies with $(B/T)_{fn}$ $>$ 0.75 do not require a bulge+disk model to fit their light profiles. |
If we set a value of Pig<0.32 as the (1o) threshold below which ealaxies are likely to be genuine bulge|disk svstems. then the fraction of these galaxies in our cutive SDSS sample is 26%. | If we set a value of $P_{pS} \leq 0.32$ as the $\sigma$ ) threshold below which galaxies are likely to be genuine bulge+disk systems, then the fraction of these galaxies in our entire SDSS sample is $\%$. |
The quality of the SDSS imagine is iusufficieut to determine bulge Sérrsic iudices for galaxies in our selected range of apparent magnitudes as 77, versus bulee fraction does notshow any statistically sijenificaut differences between a,= Laud free à, models. | The quality of the SDSS imaging is insufficient to determine bulge Sérrsic indices for galaxies in our selected range of apparent magnitudes as $P_{n4}$ versus bulge fraction does notshow any statistically significant differences between $n_b=4$ and free $n_b$ models. |
Oulv 9% of the galaxies have DP,,<0.32. | Only $\%$ of the galaxies have $P_{n4} \leq 0.32$. |
Figue 1 illustrates the usefulness of the F-test Xobabilities to select genuine bulge|disk systems. | Figure \ref{ftest_inc.ps} illustrates the usefulness of the $F$ -test probabilities to select genuine bulge+disk systems. |
The axial ratio distribution of a sample of disks raucdomly inclined iu space should be uniform between zero and oue inodulo some perturbations due to dust and/or xus. | The axial ratio distribution of a sample of disks randomly inclined in space should be uniform between zero and one modulo some perturbations due to dust and/or bars. |
Looking at the top row of Figure 13.. one cau see hat the observed distribution is indeed uniformi at low lee fraction. but that it also tends towards the same axial ratio distribution as bulges towards higher bulec ractions if no F-test selection is made. | Looking at the top row of Figure \ref{ftest_inc.ps}, one can see that the observed distribution is indeed uniform at low bulge fraction, but that it also tends towards the same axial ratio distribution as bulges towards higher bulge fractions if no $F$ -test selection is made. |
This behavior would argue that “disks” iu highly bulge-doiminated ealaxies are not in fact real but were rather introduced by he fitting algorithia as au additional deeree of freedom o 1odol the outer wines of a single courponeut galaxy. | This behavior would argue that “disks" in highly bulge-dominated galaxies are not in fact real but were rather introduced by the fitting algorithm as an additional degree of freedom to model the outer wings of a single component galaxy. |
If we select oulv galaxies with Posx:0.32. then one cau see that the resulting disk axial ratio remains uniforiu even at the highest bulge fractious. | If we select only galaxies with $P_{pS} \leq 0.32$, then one can see that the resulting disk axial ratio remains uniform even at the highest bulge fractions. |
We are thus able to select real disks in highly bulee dominated ealaxies. | We are thus able to select real disks in highly bulge dominated galaxies. |
A detailed comparison between our three different fitting models could casily παπα on its own as a separate oper. so we brieflY comment here on the salicut points of this comparison. | A detailed comparison between our three different fitting models could easily stand on its own as a separate paper, so we briefly comment here on the salient points of this comparison. |
We start by looking at 1icasurenments of the global galaxy structure from the different models iunely the galaxy Sérrsic iudex vy aud the bulge fraction BiT. | We start by looking at measurements of the global galaxy structure from the different models namely the galaxy Sérrsic index $n_g$ and the bulge fraction $B/T$ . |
The distribution of galaxy Sérrsic iudex values iu Figure LL shows a peak a sharp peak at os, = 0.5. a woader and larger peak at on4 l. no peak at n, = Lauda peak at the maxima allowed Sévrsic value of 8. | The distribution of galaxy Sérrsic index values in Figure \ref{dr7_fn_n4_pS_cmp_struct} shows a peak a sharp peak at $n_g$ = 0.5, a broader and larger peak at $n_g$ = 1, no peak at $n_g$ = 4 and a peak at the maximum allowed Sérrsic value of 8. |
Wo will return to the peaks at vg = 0.5 aud vy = 5 ater. | We will return to the peaks at $n_g$ = 0.5 and $n_g$ = 8 later. |
The large peak at i, = 1 refiects the fact that the οσα]. galaxy population is dominated bv disk ealaxics. | The large peak at $n_g$ = 1 reflects the fact that the local galaxy population is dominated by disk galaxies. |
The lack of a stroug peak at ag=Lis quite interesting )ecause it arenes for a lack of global structural similarity in galaxies that are not disk-dominuated. | The lack of a strong peak at $n_g = 4$ is quite interesting because it argues for a lack of global structural similarity in galaxies that are not disk-dominated. |
In terius of bulec action. measured values frou both the 0;=Laud free my fits (denoted as (B/T),, aud (B/T)y,, respectively verve} are well correlated with one auother aud with à», (Figure 11). | In terms of bulge fraction, measured values from both the $n_b = 4$ and free $n_b$ fits (denoted as $(B/T)_{n4}$ and $(B/T)_{fn}$ respectively here) are well correlated with one another and with $n_g$ (Figure \ref{dr7_fn_n4_pS_cmp_struct}) ). |
(B/T),; aud (B/T)g, increase with ny aud reach a value of oue at vy=L | $(B/T)_{n4}$ and $(B/T)_{fn}$ increase with $n_g$ and reach a value of one at $n_g = 4$. |
Two “brauches” are seen bevoud ay=LE: one at B/T~ I and another oue at B/T~0.5. | Two “branches" are seen beyond $n_g = 4$: one at $B/T \sim$ 1 and another one at $B/T \sim 0.5$. |
The first brauch is not suprising. but the second one may point to a potential sinele- versus double-coniponeut degeneracy in the bulee|disk decompositions. | The first branch is not surprising, but the second one may point to a potential single- versus double-component degeneracy in the bulge+disk decompositions. |
Iu order to uuderstaud the behavior of (B/T) at high πμ. we looked at the bulge aud disk ellipticity difference Aeστον64 versus the ratio of the bulee and disk ΠαΠο radii Aye,—r./(1.67r4) in the range 7€n4<δ for two rvauges in (DT), Olx(BT),0.7 aud 0.9x(οΤὸ«1.0. | In order to understand the behavior of $(B/T)$ at high $n_g$, we looked at the bulge and disk ellipticity difference $\Delta e \equiv e_b - e_d$ versus the ratio of the bulge and disk half-light radii $R_{b/d} \equiv r_e/(1.67r_d)$ in the range $7 \leq n_g <8$ for two ranges in $(B/T)_{n4}$: $0.4 \leq (B/T)_{n4}\leq 0.7$ and $0.9 \leq (B/T)_{n4}< 1.0$. |
Ae is zero Which means| that the bulges aud disks have the sale ellipticities be.. these galaxies may actually have a single compoucut but the fitting aleorithm may be using a bulee and disk components to model something that it cannot do even with (5/T),, = 1 for nj = IL. | $\Delta e$ is zero which means that the bulges and disks have the same ellipticities i.e., these galaxies may actually have a single component but the fitting algorithm may be using a bulge and disk components to model something that it cannot do even with $(B/T)_{n4}$ = 1 for $n_b$ = 4. |
The dichotomy actually comes from the ratio of the radi. | The dichotomy actually comes from the ratio of the radii. |
Galaxies with 0.1xCG/T),,<0.7 all have A5; values of 0.13 with a dispersion of 0.05 whereas galaxies with 49«(GD/T),,<1.0 have values around 0.56 with a dispersion of 0.16. | Galaxies with $0.4 \leq (B/T)_{n4}\leq 0.7$ all have $R_{b/d}$ values of 0.13 with a dispersion of 0.05 whereas galaxies with $0.9 \leq (B/T)_{n4} < 1.0$ have values around 0.56 with a dispersion of 0.16. |
There are six times more galaxies with LE(B/T),|L7 than with 0.9«(B/T),,1.0. | There are six times more galaxies with $0.4 \leq (B/T)_{n4} \leq 0.7$ than with $0.9 \leq (B/T)_{n4} < 1.0$. |
The peak in Rye, where a galaxy oeuds up soeurs to depend ou the spread in Ae. | The peak in $R_{b/d}$ where a galaxy ends up seems to depend on the spread in $\Delta e$. |
The dispersion in Ac or galaxies with 0.1<(B/T),,O.7 is twice the dispersion in Ae for galaxies with 0.9<(οΤὸ« 1.0. | The dispersion in $\Delta e$ for galaxies with $0.4 \leq (B/T)_{n4} \leq 0.7$ is twice the dispersion in $\Delta e$ for galaxies with $0.9 \leq (B/T)_{n4} < 1.0$ . |
A smaller Ae value for a given galaxy makesit more ikelv that the fitting aleovitlin will tend towards a suele compoucut model rather than a two-component uodel because the algoritlin will need two componcuts o inodel a change of ellipticity with radius that does rot come from PSF simearing. | A smaller $\Delta e$ value for a given galaxy makesit more likely that the fitting algorithm will tend towards a single component model rather than a two-component model because the algorithm will need two components to model a change of ellipticity with radius that does not come from PSF smearing. |
There is no siguificaut | There is no significant |
quality, but they demonstrate the same properties (as far as the statistics allow) of power spectra of their variability and cross-corelation functions. | quality, but they demonstrate the same properties (as far as the statistics allow) of power spectra of their variability and cross-corelation functions. |
Authors thank Coel Hellier for useful comments about accretion curtains in intermediate polars. | Authors thank Coel Hellier for useful comments about accretion curtains in intermediate polars. |
This research made use of data obtained from the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard Space Flight Center. | This research made use of data obtained from the High Energy Astrophysics Science Archive Research Center Online Service, provided by the NASA/Goddard Space Flight Center. |
This work was supported by a grant of Russian Foundation of Basic Research 10-02-00492-a, NSh-5069.2010.2, and program of Presidium of RAS "The origin and evolution of stars and galaxies" (P-19). | This work was supported by a grant of Russian Foundation of Basic Research 10-02-00492-a, NSh-5069.2010.2, and program of Presidium of RAS “The origin and evolution of stars and galaxies” (P-19). |
We concern ourselves with the general dynamical evolution of a self-gravitating stellar svstem dominated. to the first-order bv a. stable. or. slowly evolving spheriodal mass distribution but with significant and dynamically distinct substructure. | We concern ourselves with the general dynamical evolution of a self-gravitating stellar system dominated to the first-order by a stable or slowly evolving spheriodal mass distribution but with significant and dynamically distinct substructure. |
We can refer to this as an. inhomogencous spheroidal system. a disk or other structure embedded within a dark matter halo. an encounter between an elliptical galaxy and less-massive companion. sinking satellites. rings. and fine structure in ellipticals. | We can refer to this as an inhomogeneous spheroidal system, a disk or other structure embedded within a dark matter halo, an encounter between an elliptical galaxy and less-massive companion, sinking satellites, rings, and fine structure in ellipticals. |
The aim in this work is to develop an cllictent and. handy method of modelling the complex cvnamical evolution of the types of systems mentioned. given the resources which are commonly at hand. | The aim in this work is to develop an efficient and handy method of modelling the complex dynamical evolution of the types of systems mentioned, given the resources which are commonly at hand. |
An established method of investigating the dynamical evolution of stellar svstems is through the use of N-body simulations. | An established method of investigating the dynamical evolution of stellar systems is through the use of N-body simulations. |
Such simulations have become increasinglv sophisticated over. the last decade. enabling detailed experiments to be performed. on N-bodsy models of stellar systems. | Such simulations have become increasingly sophisticated over the last decade, enabling detailed experiments to be performed on N-body models of stellar systems. |
Improvements have been a result’ of greater computing power and more cunning algorithms for following 1 time evolution of à system of particles. | Improvements have been a result of greater computing power and more cunning algorithms for following the time evolution of a system of particles. |
X major 'onsideration for the researcher is to match the computing resources available to an N-body method suited to modelling 10 svstem under consideration. | A major consideration for the researcher is to match the computing resources available to an N-body method suited to modelling the system under consideration. |
The majority of numerical treatments of. stellar lvnamies rely upon the assumption that stellar svstems on 1e scale of galaxies are collision free. | The majority of numerical treatments of stellar dynamics rely upon the assumption that stellar systems on the scale of galaxies are collision free. |
This is based on the fact that the two-bocky relaxation time of a star. Fidi Is many magnitudes larger than the age of the galaxy which contains it. | This is based on the fact that the two-body relaxation time of a star, $t_{\rm relax}$, is many magnitudes larger than the age of the galaxy which contains it. |
The appropriate globally averaged estimate is eiven by.lovoss- for a svstem ol N bodies where the crossing time. foro... is the time for a particleto cross the svstem once (BinneyTremaine 1987). | The appropriate globally averaged estimate is given by,, for a system of $N$ bodies where the crossing time, $t_{\rm cross}$, is the time for a particleto cross the system once \cite{bt}. |
. Phe two-bocly relaxation rate of a body in) self-eravitating svstems depends in part on the local density aud velocity. dispersion. | The two-body relaxation rate of a body in self-gravitating systems depends in part on the local density and velocity dispersion. |
Relaxation of the orbits of individual particles will occur more quickly in regions of higher density or lower νε‘locity dispersion. | Relaxation of the orbits of individual particles will occur more quickly in regions of higher density or lower velocity dispersion. |
Other authors present more detailed. discussions. of relaxation processes. J|arouki Salpeter 1994: Luang 1993). | Other authors present more detailed discussions of relaxation processes Farouki Salpeter 1994; Huang 1993). |
LE one. is concentrating on short timescale evolution. in à region of fine substructure in an otherwise dynamically stable or only slowly evolving svstem. one may find that the resolution locally is insullicient to accurately describe the detailed evolution. | If one is concentrating on short timescale evolution in a region of fine substructure in an otherwise dynamically stable or only slowly evolving system, one may find that the resolution locally is insufficient to accurately describe the detailed evolution. |
“Phis provides the motivation to combine simulation techniques which will deal. seperately but. elliciently with cifferent components of an N-body svsteor. | This provides the motivation to combine simulation techniques which will deal seperately but efficiently with different components of an N-body system. |
The dynamical evolution of a svstem of collisionless particles is described by the collisionless Boltzmann equation. | The dynamical evolution of a system of collisionless particles is described by the collisionless Boltzmann equation, f |
our model for the case of ceutered multi-scale moments. | our model for the case of centered multi-scale moments. |
To do this. all quantities Ny;- are centered on the same ≼∐↥⋅↸∖≼⊳↑↕∪∐∪∐↑∐↸∖↴∖↴↘↽∙↖↽↴⋝∏↑↖↖⇁↸∖⋜↧∐∪↖↖⇁↑∐↸∖⋜⋯∶↴∙⊾∏↕⋜∐⋅ to ∙⋅ different. | To do this, all quantities $X_{si}$ are centered on the same direction on the sky but we allow the angular radii $\theta_{si}$ to be different. |
Then.» Eqs.(21)) aud (11)) Moll generalize as be! d (51) (PG) and Ee 03 ↽ | Then, \ref{Xs-variance}) ) and \ref{Xs-3-real}) ) generalize as = ) and = _1) _2) _3) _2 _3 _3). |
∖↽⊽↾"NUM fο lig | We show in Fig. |
We show in Fig. 7 the two-scale secoud-order womeuts ο and May(Mau(o0)? for a scale-ratio a=2.5. and 10. at τν=1. | \ref{fig_alpha_2} the two-scale second-order moments $\lag \kappa_s(\theta_s)\kappa_s(\alpha\theta_s)\rag$ and $\lag \Map(\theta_s)\Map(\alpha\theta_s)\rag$ for a scale-ratio $\alpha=2, 5$, and $10$, at $z_s=1$. |
A higher a vields a stnaller moment because it corresponds to a larger second aneular radius οος, | A higher $\alpha$ yields a smaller moment because it corresponds to a larger second angular radius $\alpha\theta_s$. |
We obtain the same level of agreement as for the simele-scale variances shown iu Fie. 1l. | We obtain the same level of agreement as for the single-scale variances shown in Fig. \ref{fig_xi}. |
In particular. we obtain a good match on small angular scales where the “halo-fit™ foiuula somewhat undoerestinates the weak-leusiug power. | In particular, we obtain a good match on small angular scales where the “halo-fit” formula somewhat underestimates the weak-lensing power. |
On large scales our analytical results are somewhat larecr than the ata obtained from the munerical simulations. | On large scales our analytical results are somewhat larger than the data obtained from the numerical simulations. |
This is due to the missing of larec- modes in the simulations because of the fuite size | This is due to the missing of large-scale modes in the simulations because of the finite size |
host galaxy of the quasar as marginally more extended than the (combined. MOS|ARGUS) instrumental point-spread function (PSE). similar to the (slightly asymmetric) distribution of the quasar in seen in broadLh. | host galaxy of the quasar as marginally more extended than the (combined MOS+ARGUS) instrumental point-spread function (PSF), similar to the (slightly asymmetric) distribution of the quasar in seen in broad. |
We present reconstructed. images of 3€323.1 in the light of broad. and OLI] in Fig 1: a disc is drawn at the position of each fibre where line emission was significantly. detected areas left blank within the outline of the ARGUS aperture have no line detection in those fibres. | We present reconstructed images of 3C323.1 in the light of broad and [OIII] in Fig \ref{fig:3c323.1im}: a disc is drawn at the position of each fibre where line emission was significantly detected – areas left blank within the outline of the ARGUS aperture have no line detection in those fibres. |
The grevscale colour within each disk gives the intensity of the line emission measured within that fibre’s spectrum. | The greyscale colour within each disk gives the intensity of the line emission measured within that fibre's spectrum. |
Vig 1 shows the OL] line emission to be clearly more extended than he approximate SP given by the broad image. | Fig \ref{fig:3c323.1im} shows the [OIII] line emission to be clearly more extended than the approximate PSF given by the broad image. |
Phe nebula extends out o à maximum of 4.1 arcsec kpe)) to the East (at p.a. 105). | The nebula extends out to a maximum of 4.1 arcsec ) to the East (at p.a. ), |
and to 1.9 arcsec (~LOkpe) in all other directions. with the maximum cimension approximately perpendicular o the radio source axis. ( | and to $\sim1.9$ arcsec $\sim$ 10kpc) in all other directions, with the maximum dimension approximately perpendicular to the radio source axis. ( |
1 aresee corresponds to ~5.1 aat the redshift of the quasar. assuming the cosmology of 44,—250kms tand go=0.5 whieh will be used hroughout this paper) | 1 arcsec corresponds to $\sim5.1$ at the redshift of the quasar, assuming the cosmology of $H_0$ and $q_0$ =0.5 which will be used throughout this paper). |
Redshifted line emission is barely detected in the IELIt. even when the spectra are binned into larger T-libre eells. as it falls in the less sensitive region of the combined | grism response. | Redshifted line emission is barely detected in the EELR, even when the spectra are binned into larger 7-fibre cells, as it falls in the less sensitive region of the combined $+$ grism response. |
Narrow is also too faint o be significantly detected outside of the nuclear regions in individual fibres. | Narrow is also too faint to be significantly detected outside of the nuclear regions in individual fibres. |
We obtain an oll-nuclear detection of each of OL] and only by summing the spectra [rom 20 Libres orming the most extended region at the SIS of the nebula (the fibres marked by diamonds in Fig 2)). | We obtain an off-nuclear detection of each of [OII] and only by summing the spectra from 20 fibres forming the most extended region at the SE of the nebula (the fibres marked by diamonds in Fig \ref{fig:3c323.1diag}) ). |
Fitting the total spectrum from this region enables us to obtain average line luxes (although OLI] had to be fixed at the same velocity width as the OH] anc complex) and thus intensity ratios of and OL]/OLLI] (Lable 2 and Fig 4)). | Fitting the total spectrum from this region enables us to obtain average line fluxes (although [OII] had to be fixed at the same velocity width as the [OIII] and complex) and thus intensity ratios of and [OII]/[OIII] (Table \ref{tab:linerats} and Fig \ref{fig:allrats}) ). |
We caleulate the fractions of total ine emission emitted [rom the nucleus (defined. as emitted by the central. 19. fibres) in the both OLLI] and OLL and find them to be 69 and 52 per cent respectively (Table 33). | We calculate the fractions of total line emission emitted from the nucleus (defined as emitted by the central 19 fibres) in the both [OIII] and [OII] and find them to be 69 and 52 per cent respectively (Table \ref{tab:nebprops}) ). |
“Phese may well represent. underestimates of the actual fractions. as we do not take into account any EELR hat may be contributingIn emission alonge the line of sighte to he quasar nucleus. | These may well represent underestimates of the actual fractions, as we do not take into account any EELR that may be contributing emission along the line of sight to the quasar nucleus. |
Fig 1. also shows the radial velocity of the ine relative to that the peak fibre. | Fig \ref{fig:3c323.1im} also shows the radial velocity of the line relative to that the peak fibre. |
The nebula appears o show a fairly regular cipolar motion with amplitude of E200. albeit about an axis misaligned from that of he radio source by approximately 25- degrees. | The nebula appears to show a fairly regular dipolar motion with amplitude of $\pm$, albeit about an axis misaligned from that of the radio source by approximately 25 degrees. |
The radial velocity values along cuts at position angles of907.. and (as marked in Fig 2)) are shown in Fig 3.. | The radial velocity values along cuts at position angles of, and (as marked in Fig \ref{fig:3c323.1diag}) ) are shown in Fig \ref{fig:3c323.1rvint}. |
The linewidth (FWIIM as measured [rom the fit to the OLLI] lines). is righest around the nucleus. dropping to 200400 in he extended emission to the SE. | The linewidth (FWHM as measured from the fit to the [OIII] lines) is highest around the nucleus, dropping to 200–400 in the extended emission to the SE. |
4C11.72 is a relatively compact (diameter τὸ kpce)). linear radio source associated with a low-redshift quasar located in a small cluster of galaxies (Gunn 1971: Robinson Wampler 1972: Yee Green 1984: Block Stockton 1991: Ellingson 19904). | 4C11.72 is a relatively compact (diameter $\sim75$ ), linear radio source associated with a low-redshift quasar located in a small cluster of galaxies (Gunn 1971; Robinson Wampler 1972; Yee Green 1984; Block Stockton 1991; Ellingson 1994). |
Phe quasar lies in a relatively uncisturbed host galaxy (Llutchings Nell 1992). but has an OLLI] emission-line region extended. alone position angle1003. with an orientation ancl a spatial scale comparable to the radio source (SAIST: Llutchines Crampton 1990: Durret 1994). | The quasar lies in a relatively undisturbed host galaxy (Hutchings Neff 1992), but has an [OIII] emission-line region extended along position angle, with an orientation and a spatial scale comparable to the radio source (SM87; Hutchings Crampton 1990; Durret 1994). |
4C11.72 is the best case amongst. low-redshilt quasars for an association between the EIZLIU and. racio emission. although not all the ionizecl gas is related. to the | 4C11.72 is the best case amongst low-redshift quasars for an association between the EELR and radio emission, although not all the ionized gas is related to the |
implied by ((7)) might thus only hold for the solar neighborhood. | implied by \ref{eq:mass-reference_large-radii}) ) might thus only hold for the solar neighborhood. |
It is obvious that the formation of a cluster requires a larger mass reservoir than necessary to form a single isolated star. | It is obvious that the formation of a cluster requires a larger mass reservoir than necessary to form a single isolated star. |
Thus, one might naively expect that regions containing clusters are more massive than those devoid of stellar groups. | Thus, one might naively expect that regions containing clusters are more massive than those devoid of stellar groups. |
If true, one would thus expect that, within a given cluster-forming cloud, the regions containing clusters are more massive than cluster-less regions of similar size. | If true, one would thus expect that, within a given cluster-forming cloud, the regions containing clusters are more massive than cluster-less regions of similar size. |
This hypothesis is tested in refsec:clusters-dominate-host.. | This hypothesis is tested in \\ref{sec:clusters-dominate-host}. |
Also, one would expect that cluster-forming cloud fragments are more massive than all similar-sized fragments in clouds not containing clusters. | Also, one would expect that cluster-forming cloud fragments are more massive than all similar-sized fragments in clouds not containing clusters. |
Section 3.3.0 investigates this issue. | Section \ref{sec:clusters-dominate-isolated} investigates this issue. |
In our sample, the Pipe Nebula and Taurus serve as examples of regions dominated by isolated star formation. | In our sample, the Pipe Nebula and Taurus serve as examples of regions dominated by isolated star formation. |
Actually, except for B59, the Pipe Nebula does hardly form stars at all | Actually, except for B59, the Pipe Nebula does hardly form stars at all \citep{forbrich2009:pipe-yso}. |
Perseus and Ophiuchus serve as examples for cluster-forming(?).. regions. | Perseus and Ophiuchus serve as examples for cluster-forming regions. |
They do contain clusters much more significant than any stellar aggregate found in Taurus and the PipeNebula?. | They do contain clusters much more significant than any stellar aggregate found in Taurus and the Pipe. |
. Orion A is another example of a cluster-forming cloud. | Orion A is another example of a cluster-forming cloud. |
We examine whether cluster-forming cloud fragments dominate the mass reservoir of their host cloud at all radii. | We examine whether cluster-forming cloud fragments dominate the mass reservoir of their host cloud at all radii. |
This is executed in reffig:cloud-sample.. | This is executed in \\ref{fig:cloud-sample}. |
Here, we only consider cluster-forming clouds with high quality data, PPerseus and Ophiuchus. | Here, we only consider cluster-forming clouds with high quality data, Perseus and Ophiuchus. |
Inspection of the column density maps reveals that the most massive small-scale features in Perseus and (highlighted)Ophiuchus are located in the NGC1333 and L1688 clusters, respectively. | Inspection of the column density maps reveals that the (highlighted) most massive small-scale features in Perseus and Ophiuchus are located in the NGC1333 and L1688 clusters, respectively. |
At small spatial scales, the most massive fragments are thus indeed located in cluster-forming regions. | At small spatial scales, the most massive fragments are thus indeed located in cluster-forming regions. |
Note, though, that cluster-forming regions do also contain fragments of lower mass. | Note, though, that cluster-forming regions do also contain fragments of lower mass. |
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