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The errors listed iu both Table 23 and Lave directly from the SPECFIT task. which assuues that the errors on the iuput spectrum follow a Cassia distribution. | The errors listed in both Table 3 and 4 are directly from the SPECFIT task, which assumes that the errors on the input spectrum follow a Gaussian distribution. |
The tabulated errors represcut oue sigma for asinele imteresting parameter. | The tabulated errors represent one sigma for a single interesting parameter. |
We fit a PL οὐαμα using these relatively line-free regions: A LO20-1050. (150-1270. £120-L150. 1210-E780. and 5070-5130. | We fit a PL continuum using these relatively line-free regions: $\lambda$ 4020-4050, 4150-4270, 4420-4450, 4710-4780, and 5070-5130. |
The resulting fit parameters are fixed. while the following single Gaussian components are fit to the strong emission lines between 1150 to5100À:: 13 BroadIT: 2) Narrowο, with EWITIM fixed to that of A5007: 3) A5007: 0) A1959. fixed relative to the central wavelength aud (1/3) fiux of component (3): Finally. the blended iron cussion above the PL continu is stutiedd over the wavelength rauge AALISL£681. | The resulting fit parameters are fixed, while the following single Gaussian components are fit to the strong emission lines between 4150 to: 1) Broad; 2) Narrow, with FWHM fixed to that of $\,\lambda$ 5007; 3) $\,\lambda$ 5007; 4) $\,\lambda$ 4959, fixed relative to the central wavelength and (1/3) flux of component (3); Finally, the blended iron emission above the PL continuum is summed over the wavelength range $\lambda\lambda 4434-4684$. |
As can be seen from Figure 3 and Table 3. by far the strongest effect is the great relative strength of eniission in the NB composite: Cluission is 2.5 times stronger in X-ray bright sample. | As can be seen from Figure 3 and Table 3, by far the strongest effect is the great relative strength of emission in the XB composite; emission is $\sim 2.5$ times stronger in X-ray bright sample. |
Broad chussiou is about strouger m the NB composite. although the FWIDI are similar at ~5700 | Broad emission is about stronger in the XB composite, although the FWHM are similar at $\sim$. |
Streneths of the narrow component of are comparable. | Strengths of the narrow component of are comparable. |
Ta both NB aud NF composites. broad is redshifted relative to rest (1..6.. | In both XB and XF composites, broad is redshifted relative to rest .e., |
relative to A5007 or narrow ΠΠ.) bv aboutl | relative to $\,\lambda$ 5007 or narrow ) by about. |
o Doroson Creen (1992: hereafter BO92). mi a principal colmponent analysis of low redshift QSO spectra. found that the principal cigenvector of their sample (the lear colubination of parameters that represents the largest | Boroson Green (1992; hereafter BG92), in a principal component analysis of low redshift QSO spectra, found that the principal eigenvector of their sample (the linear combination of parameters that represents the largest |
canonical NSs lave faded. | canonical NSs have faded. |
Detection of spin-period variation in N-raw flux over several N-vay observations wold explore periods »10! sec for the case ofJ1856. | Detection of spin-period variation in X-ray flux over several X-ray observations would explore periods $>10^4$ sec for the case of. |
. Ta most cases in which has spuu-down to P>LO! sec (83)). it should at preseut have a period of at least 109 sec. | In most cases in which has spun-down to $P>10^4$ sec \ref{sec_1856}) ), it should at present have a period of at least $10^6$ sec. |
The chance of κοπο a different surface aspect of the star may be small and eravitational Ποιοιπιο effects will siuear out some spectral modulation. | The chance of seeing a different surface aspect of the star may be small and gravitational light-bending effects will smear out some spectral modulation. |
Tf an isolated NS in a propeller phase is confirmed it will select amone propeller effect models. | If an isolated NS in a propeller phase is confirmed it will select among propeller effect models. |
Ifa NS is in the propeller phase. oue should be cautious in deriving NS ages and magnetic field streneths frou: 2? aud P'ueasuremeuts since the conveutional dipole radiation wind foxiulae of equation (1)) aud (2)) are no longer valid. | If a NS is in the propeller phase, one should be cautious in deriving NS ages and magnetic field strengths from $P$ and $\dot{P}$ measurements since the conventional dipole radiation wind formulae of equation \ref{eq_canonical}) ) and \ref{eq_csd1}) ) are no longer valid. |
In such cases. independent determination of ages aud magnetic field stronethls (οιο, supernova remnant ages aud spectroscopic measurements of B-ficld) ave needed. | In such cases, independent determination of ages and magnetic field strengths (e.g. supernova remnant ages and spectroscopic measurements of B-field) are needed. |
that convergence is only achieved when the cooling lengths are aclequately resolved. | that convergence is only achieved when the cooling lengths are adequately resolved. |
At the moment. however. it is not feasible to perform simulations of a shock-cloud interaction in 3D that attain the required. resolution. | At the moment, however, it is not feasible to perform simulations of a shock-cloud interaction in 3D that attain the required resolution. |
Therefore. we adopt a resolution similar to that used. by Shin.Stone.&Suvder(2008). for the adiabatie shocks. | Therefore, we adopt a resolution similar to that used by \cite{SSS08} for the adiabatic shocks. |
We use an adaptive mesh with 5 levels of refinement with the resolution at the finest grid 4807 608. | We use an adaptive mesh with 5 levels of refinement with the resolution at the finest grid $^2 \times$ 608. |
Such a resolution corresponds to a physical eric spacing of 1.58 pe. | Such a resolution corresponds to a physical grid spacing of 1.58 pc. |
As translucent clumps in GMCSs have length scales of about 3 pe. it is clear that we cannot resolve such structures. | As translucent clumps in GMCs have length scales of about 3 pc, it is clear that we cannot resolve such structures. |
Our simulations. thus. can only show the onset of clump formation. | Our simulations, thus, can only show the onset of clump formation. |
As the cloud is accelerated to move with the post-shock llow (e.g.MacLowetal.1994).. the cloucl eventually moves olf the grid. | As the cloud is accelerated to move with the post-shock flow \citep[e.g.][]{Metal94}, the cloud eventually moves off the grid. |
To avoid this we calculate the densitv-weighed average velocity of the cloud. along each dillerent. axis at every timestep. | To avoid this we calculate the density-weighed average velocity of the cloud along each different axis at every timestep. |
The densitv-weighed average of cach velocity component is defined as where f is ορ 04 or ον and C'ds a scalar which is 1 for cloud material and 0 for ambient gas. | The density-weighed average of each velocity component is defined as where $f$ is $v_x$, $v_y$ or $v_z$ and $C$ is a scalar which is 1 for cloud material and 0 for ambient gas. |
p.V. ancl Ady are the mass density. volume. and cloud mass. | $\rho, V$ and $M_{cl}$ are the mass density, volume, and cloud mass. |
By performing a CGalilean. transformation of the computational domain. the cloud. remains in the centre of the grid. | By performing a Galilean transformation of the computational domain, the cloud remains in the centre of the grid. |
Furthermore. we are able to study the acceleration of the cloud. | Furthermore, we are able to study the acceleration of the cloud. |
Note that densitv-weighed. averages can also be calculated for other How variables. | Note that density-weighed averages can also be calculated for other flow variables. |
The shock-cloud interaction for a parallel. shock in 3D is nearly identical to the evolution studied in the 2D axisvmmetrie simulations of Paper lL. Lt is useful to describe he dvnamical evolution here again. not only to show the dillerences compared. to the 2D. results. but. also. because some of the dynamical characteristics are relevant for the »erpendieular and oblique models. | The shock-cloud interaction for a parallel shock in 3D is nearly identical to the evolution studied in the 2D axisymmetric simulations of Paper I. It is useful to describe the dynamical evolution here again, not only to show the differences compared to the 2D results, but also because some of the dynamical characteristics are relevant for the perpendicular and oblique models. |
Figure 1 shows results for the parallel case at two imes. | Figure \ref{fig:paboundary} shows results for the parallel case at two times. |
As the intercloud shock sweeps around the cloud (in a time of /,4= 24/0). à shock is transmitted into the cloud ancl a bow-shock forms in. front of the cloud. | As the intercloud shock sweeps around the cloud (in a time of $t_{cp} = 2R_{cl}/v_{ext}$ ), a shock is transmitted into the cloud and a bow-shock forms in front of the cloud. |
The ransmuttect fast-mode shock has a lower propagation speed han the intercloud shock. Le. tiny—couNE? where \ is the density ratio of cloud/intercloud. gas. | The transmitted fast-mode shock has a lower propagation speed than the intercloud shock, i.e. $v_{int} = v_{ext}/\chi^{1/2}$ where $\chi$ is the density ratio of cloud/intercloud gas. |
Pherefore. a velocity shear laver forms at the cloud boundary. | Therefore, a velocity shear layer forms at the cloud boundary. |
Because of this slip surface. a vortex ring develops and sweeps cloud material away from the cloud. (see Fig. 1)). | Because of this slip surface, a vortex ring develops and sweeps cloud material away from the cloud (see Fig. \ref{fig:paboundary}) ). |
As the fast-niocle shock moves through the cloud. it compresses and heats the initially thermallv-stable warm eas so that it ends up being thermally unstable. | As the fast-mode shock moves through the cloud, it compresses and heats the initially thermally-stable warm gas so that it ends up being thermally unstable. |
Figure. | Figure. |
2 shows the distribution of mass in thepon phase space at two different times. | \ref{fig:pa_phase} shows the distribution of mass in the $p-n$ phase space at two different times. |
Phe time scale for radiative cooling is shorter than the propagation time of the fast-mode shock through the cloud. given by the cloud-crushing time ἐν=iyferns. | The time scale for radiative cooling is shorter than the propagation time of the fast-mode shock through the cloud, given by the cloud-crushing time $t_{cc} = R_{cl}/v_{int}$. |
Thus. the gas loses a significant. fraction. of its internal energy during the compression. | Thus, the gas loses a significant fraction of its internal energy during the compression. |
Furthermore. the magnetic pressure increases behind the fast-mode shock. | Furthermore, the magnetic pressure increases behind the fast-mode shock. |
Hence. the value of 3 drops below unity inside the cloud. | Hence, the value of $\beta$ drops below unity inside the cloud. |
Εις provides the ideal conditions for the generation. by MILD. waves. of dense clumps and cores (Falle&Lartquist2002:VanLoo.Falle.&Llartquist2006:VanLooetal. 2008). | This provides the ideal conditions for the generation, by MHD waves, of dense clumps and cores \citep{FH02, VFH06, VFH08}. |
. C'old. dense clumps can also form clue to small perturbations along the unstable part of the equilibrium curve (Inutsuka&Ixovama2007). | Cold, dense clumps can also form due to small perturbations along the unstable part of the equilibrium curve \citep{KI06}. |
. In our simulation we find a small fraction of cloud material on the thermally unstable part of the equilibrium curve after zz 6 Myr (or 0.764). | In our simulation we find a small fraction of cloud material on the thermally unstable part of the equilibrium curve after $\approx$ 6 Myr (or $_{cc}$ ). |
Evpical timescales for both oocesses are a few Myr. | Typical timescales for both processes are a few Myr. |
As the cloud flattens and fragments in about 1.5-2 t4. (see Paper D). there is ample time for hese processes to work. | As the cloud flattens and fragments in about 1.5-2 $_{cc}$ (see Paper I), there is ample time for these processes to work. |
However. the numoerical resolution of our simulations is insullicient to establish whether these oocesses are the dominant formation mechanisms of dense clumps ancl cores within clouds. | However, the numerical resolution of our simulations is insufficient to establish whether these processes are the dominant formation mechanisms of dense clumps and cores within clouds. |
As we cannot follow this ormation process. we stop the simulation shortly. after £,,. | As we cannot follow this formation process, we stop the simulation shortly after $t_{cc}$. |
This timescale was chosen because the re-expansion phase of he cloud then starts. | This timescale was chosen because the re-expansion phase of the cloud then starts. |
Also. sclberavity which is not included in these simulations becomes globally important around this ime ancl will alfect the subsequent dynamical evolution. | Also, self-gravity which is not included in these simulations becomes globally important around this time and will affect the subsequent dynamical evolution. |
The fast-moce shock is not the only shock to be ransmitted into the cloud. | The fast-mode shock is not the only shock to be transmitted into the cloud. |
A slow-mocde shock is trailing 10. [ast-mode shock. | A slow-mode shock is trailing the fast-mode shock. |
As it moves much more slowly than 1ο [ast-mocdoe shock. it remains close to the boundary of the loud. | As it moves much more slowly than the fast-mode shock, it remains close to the boundary of the cloud. |
Behind the slow-mocde shock. the magnetic pressure ecreases and there is nothing that prevents the eas from Compressing as it cools. | Behind the slow-mode shock, the magnetic pressure decreases and there is nothing that prevents the gas from compressing as it cools. |
Figure 2. clearly shows the rapid condensation due to the slow-mode shock. | Figure \ref{fig:pa_phase} clearly shows the rapid condensation due to the slow-mode shock. |
In a few Myr 1e gas behind the slow-mocde shock cools to the thermally-μαable cold phase. | In a few Myr the gas behind the slow-mode shock cools to the thermally-stable cold phase. |
For the gas behind the fast-nioce shock mit is not processed by the slow-mode shock. this process occurs much more slowly. | For the gas behind the fast-mode shock that is not processed by the slow-mode shock, this process occurs much more slowly. |
Thus. a dense. cold. laver quickly forms at the cloud boundary with the highest densities on the upstream, parts of the cloud where the shock first hits the cloud (see Fie. 1)). | Thus, a dense, cold layer quickly forms at the cloud boundary with the highest densities on the upstream parts of the cloud where the shock first hits the cloud (see Fig. \ref{fig:paboundary}) ). |
While the cooling behind the slow-mioce shock is thermally. unstable. the dense shell is subject to Ixelvin-LEelmholtz and Ravleigh-VTavlor instabilities. even though the dynamical instabilities are mostly suppressed by the strong magnetic Ποιά (c.g.Chancrasckhar1961). | While the cooling behind the slow-mode shock is thermally unstable, the dense shell is subject to Kelvin-Helmholtz and Rayleigh-Taylor instabilities, even though the dynamical instabilities are mostly suppressed by the strong magnetic field \citep[e.g.][]{C61}. |
. Hence. the shell breaks up into dense fragments. | Hence, the shell breaks up into dense fragments. |
The censest clumps have number densities of z10 ? and are à few parsee in size. | The densest clumps have number densities of $\approx 10^3$ $^{-3}$ and are a few parsec in size. |
Note that we found roughly the same values for the higher resolution 2D simulation in Paper Lo For lower resolution 2D simulations the densities are not so high. | Note that we found roughly the same values for the higher resolution 2D simulation in Paper I. For lower resolution 2D simulations the densities are not so high. |
However. the clumps of the 2D simulations are essentially. axisvmmetric rings that break up into smaller. higher density. clumps in the 3D simulations. | However, the clumps of the 2D simulations are essentially axisymmetric rings that break up into smaller, higher density clumps in the 3D simulations. |
These cold. dense clumps contain several hundreds of solar masses each and are gravitationally unstable. | These cold, dense clumps contain several hundreds of solar masses each and are gravitationally unstable. |
Thus. such clumps are likely precursors of massive stars. | Thus, such clumps are likely precursors of massive stars. |
The dynamical evolution of cloud. interacting with a perpenclicular shock is in many ways similar to that of one interacting with a parallel shock. | The dynamical evolution of cloud interacting with a perpendicular shock is in many ways similar to that of one interacting with a parallel shock. |
As the intercloucl shock sweeps around the cloud. a transmitted. fast-niocde shock propagates through the cloud. making the cloud material | As the intercloud shock sweeps around the cloud, a transmitted fast-mode shock propagates through the cloud making the cloud material |
Galaxy redshift surveys provide a cosmological probe highly complementary to the cosmic microwave background (CMB) (Penzias&Wilson1965) and supernovae (SNe) (Riessetal.1998:Perlmutteretal. 1999).. | Galaxy redshift surveys provide a cosmological probe highly complementary to the cosmic microwave background (CMB) \citep{Penzias:1965wn} and supernovae (SNe) \citep{Riess:1998cb,Perlmutter:1998np}. |
Large-scale structure data from galaxy surveys can be analyzed using either the power spectrum analysis or the correlation function analysis. | Large-scale structure data from galaxy surveys can be analyzed using either the power spectrum analysis or the correlation function analysis. |
Although these two methods are simple Fourier transforms of one another. the analysis processes are quite different and the results cannot be converted with Fourier transform directly because of the finite size of the survey volume. | Although these two methods are simple Fourier transforms of one another, the analysis processes are quite different and the results cannot be converted with Fourier transform directly because of the finite size of the survey volume. |
The SDSS data have been analyzed using both the power spectrum method (see. e.g... Tegmarketal.2004:Hutsi2005:Padmanabhan 2010). and the correlation function method (see. e.g.. etal.2009:SanchezKazin 2010). | The SDSS data have been analyzed using both the power spectrum method (see, e.g., \citealt{Tegmark:2003uf,Hutsi:2005qv,Padmanabhan:2006ku,Blake:2006kv,Percival:2007yw,Percival:2009xn,Reid:2009xm}) ), and the correlation function method (see, e.g., \citealt{Eisenstein:2005su,Okumura:2007br,Cabre:2008sz,Martinez:2008iu,Sanchez:2009jq,Kazin:2009cj}) ). |
The three major uncertainties while constructing a theoretical prediction of the power spectrum or correlation function with a given cosmological model are the galaxy bias (the relationship between galaxy and matter distributions). non-linear effects. and redshift distortions. | The three major uncertainties while constructing a theoretical prediction of the power spectrum or correlation function with a given cosmological model are the galaxy bias (the relationship between galaxy and matter distributions), non-linear effects, and redshift distortions. |
The knowledge of these uncertainties determines which analysis method and scale range we should use to obtain reliable constraints on the dark energy and cosmological parameters. | The knowledge of these uncertainties determines which analysis method and scale range we should use to obtain reliable constraints on the dark energy and cosmological parameters. |
In this paper. we present the measurement of the spherically-averaged correlation function from the SDSS DR7 luminous red galaxy (LRG) (Eisensteinetal.2001:Abazajian2009) sample which provides a homogeneous galaxy sample and has the largest effective survey volumeto date for studying the linear regime (Eisensteinetal.2005).. | In this paper, we present the measurement of the spherically-averaged correlation function from the SDSS DR7 luminous red galaxy (LRG) \citep{Eisenstein:2001cq,Abazajian:2008wr} sample which provides a homogeneous galaxy sample and has the largest effective survey volumeto date for studying the linear regime \citep{Eisenstein:2005su}. |
In Section 2.. we introduce the galaxy sample and selection functions used in this study. | In Section \ref{sec:data}, we introduce the galaxy sample and selection functions used in this study. |
In Section 3.. we describe the details of our method. | In Section \ref{sec:method}, we describe the details of our method. |
In Sec. 4.. | In Sec. \ref{sec:results}, |
we present our results. | we present our results. |
In Sec. 5.. | In Sec. \ref{sec:test}, , |
we check our results using systematic tests. | we check our results using systematic tests. |
We summarize and conclude in Sec. 6.. | We summarize and conclude in Sec. \ref{sec:conclusion}. . |
ll and K-band spectral images are shown in Figures l.. 2. and 3.. | H and K-band spectral images are shown in Figures \ref{h2}, , \ref{contours} and \ref{faint}. |
To more clearly show the line-emission features. the continuiun has been removed from each image. row by row. bv fitting a second-order polynomial to wavelength regions free of line emission. | To more clearly show the line-emission features, the continuum has been removed from each image, row by row, by fitting a second-order polynomial to wavelength regions free of line emission. |
llowever. residual shot noise associated with the bright continuum does remain in some parts of the data. | However, residual shot noise associated with the bright continuum does remain in some parts of the data. |
As is the case in other spectroscopic studies (Davisetal.2001:Takamiοἱ2005).. extended emission is observed only in II». | As is the case in other spectroscopic studies \citep{dav01,tak05}, extended emission is observed only in $_2$. |
IIowever. for the first time. faint II emission is also observed in (he counter-]et (positive ollsets in Figure 3)). | However, for the first time, faint $_2$ emission is also observed in the counter-jet (positive offsets in Figure \ref{faint}) ). |
The Brackett lines in the II-band. and the CO. aand lines in (he ας all appear as compact peaks coincident with the continuum. | The Brackett lines in the H-band, and the CO, and lines in the K-band all appear as compact peaks coincident with the continuum. |
The [Fe i] peak is also compact. although it is slightly offset along the blue-shifted jet. | The [Fe ] peak is also compact, although it is slightly offset along the blue-shifted jet. |
In Figure 4 we compare plots of the aand pprofiles traced alone the slit. together with profiles of the adjacent continuum. | In Figure \ref{profiles} we compare plots of the and profiles traced along the slit, together with profiles of the adjacent continuum. |
The continuum plots are (he average of profiles extracted slehtly blie-ward aud. rec-warcl of the line in each case. | The continuum plots are the average of profiles extracted slightly blue-ward and red-ward of the line in each case. |
The continuum proliles effectively show the spatial resolution of the observations. although they may be broadened slightly by nebulositv associated with13.. | The continuum profiles effectively show the spatial resolution of the observations, although they may be broadened slightly by nebulosity associated with. |
Similar profile plots were produced for the CO 2-0 bandheacd and lline emission (not shown). | Similar profile plots were produced for the CO 2-0 bandhead and line emission (not shown). |
Only one component is evident in each of the CO. aand [Fe ul) profiles. while as many as five separate peaks are identified in the II» profile (labeled in Figure 4)). | Only one component is evident in each of the CO, and [Fe ] profiles, while as many as five separate peaks are identified in the $_2$ profile (labeled in Figure \ref{profiles}) ). |
From Lorentzian fitting. (he offsets with respect to the source continuum position and the FWIIM of the individual components seen in each line were measured. | From Lorentzian fitting, the offsets with respect to the source continuum position and the FWHM of the individual components seen in each line were measured. |
These are given in Table 1.. | These are given in Table \ref{table1}. |
The CO and ppeaks are unresolved ancl are precisely coincident. to within a few AU. with the source continuum: we see no evidence lor CO or components associated with the outflow. | The CO and peaks are unresolved and are precisely coincident, to within a few AU, with the source continuum; we see no evidence for CO or components associated with the outflow. |
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