source
stringlengths 1
2.05k
⌀ | target
stringlengths 1
11.7k
|
|---|---|
We are interested in the streneth of these effects and whether or not thev can be distinguished. bv current or future observational data. | We are interested in the strength of these effects and whether or not they can be distinguished by current or future observational data. |
First. we consider the effect of preheating ou the temperature of a cluster. | First, we consider the effect of preheating on the temperature of a cluster. |
Figure lis a plot of Zy as a function of entropy floor level (νο, A) for three clusters of differeut.ο masses, | Figure 1 is a plot of $T_X$ as a function of entropy floor level (i.e., $K_0$ ) for three clusters of different masses. |
The thin line represeuts a cluster with AJz5.6«101A... the next thickest ue represents a cluster with AfzmLOM AL... aud the thickest line represents a cluster with AzLs.Whar,. | The thin line represents a cluster with $M \approx 5.6 \times 10^{14} M_{\odot}$, the next thickest line represents a cluster with $M \approx 10^{15} M_{\odot}$ , and the thickest line represents a cluster with $M \approx 1.8 \times 10^{15} M_{\odot}$. |
As expected. the gas temperature of a cluster mereases as the level of peheating is Increased. | As expected, the gas temperature of a cluster increases as the level of preheating is increased. |
Ou average. an increase in Py of about 1 keV (10 - i)) occurs when a cluster is preheated to the level of MoE300 keV cu? (over the range 3 keV <TyX10 keV). | On average, an increase in $T_X$ of about 1 keV (10 - ) occurs when a cluster is preheated to the level of $K_0 \gtrsim 300$ keV $^2$ (over the range 3 keV $\lesssim
T_X \lesssim 10$ keV). |
This effect will primarily manifest itself as a normalization shift iu the MaasTx relation. | This effect will primarily manifest itself as a normalization shift in the $M_{gas} - T_X$ relation. |
Figure 2 presents the dimensionless eas deusitv profile of a cluster with Fy=£L keV (eft panel) aud a cluster with Ty=8 keV (right panel) as a function of the level of preheating. | Figure 2 presents the dimensionless gas density profile of a cluster with $T_X = 4$ keV (left panel) and a cluster with $T_X = 8$ keV (right panel) as a function of the level of preheating. |
The dot-dashed Ime is the selfsimilar result) (i.c.. isothermal model of BBLP02). | The dot-dashed line is the self-similar result (i.e., isothermal model of BBLP02). |
The loue-dashed. short-dashed. dotted. aud solid lines represcut the preheated models of BBLP02 with Ay = 100. 200. 500. and £27 keV cm?. respectively. | The long-dashed, short-dashed, dotted, and solid lines represent the preheated models of BBLP02 with $K_0$ = 100, 200, 300, and 427 keV $^2$, respectively. |
Preheating reduces the gas density. aud therefore the gas mass. in the ceutral regious of a cluster. | Preheating reduces the gas density, and therefore the gas mass, in the central regions of a cluster. |
Tn we investigate the M,—Tx relation within three differentL. radii r=0.257.1 Mpe and 0.505.+ Apc aud (the radius within which the mean dark latter mass density is 500 times the mean critical deusitv Perit at 2 = O). | In 4, we investigate the $M_{gas} - T_X$ relation within three different radii: $r = 0.25 h^{-1}$ Mpc and $0.50 h^{-1}$ Mpc and $r_{500}$ (the radius within which the mean dark matter mass density is 500 times the mean critical density $\rho_{crit}$ at $z$ = 0). |
These radi are indicated in Figure 2 by the open squares. pentagous. and triangles. respective | These radii are indicated in Figure 2 by the open squares, pentagons, and triangles, respectively. |
Clearly. the effect ou the MsTy relation willbe strougest when Mj;is evaluated within r=0.25h1 Mpc. | Clearly, the effect on the $M_{gas} - T_X$ relation will be strongest when $M_{gas}$ is evaluated within $r = 0.25
h^{-1}$ Mpc. |
QO.liu Furthermore. because r=(02255+ Mpe is a fixed radius that samples different fractions of the virial radius CH,,) for clusters of different. temperature. the (fractional) reduction iu eas mass within that radius | 0.1in Furthermore, because $r = 0.25 h^{-1}$ Mpc is a fixed radius that samples different fractions of the virial radius $R_{halo}$ ) for clusters of different temperature, the (fractional) reduction in gas mass within that radius |
Recent discoveries. of exceedingly faint dwarf. spheroical galaxies have shown quite dramatically that luminosity. alone cannot be relied on to distinguish between dwarl spheroidal galaxics ancl elobular clusters. | Recent discoveries of exceedingly faint dwarf spheroidal galaxies have shown quite dramatically that luminosity alone cannot be relied on to distinguish between dwarf spheroidal galaxies and globular clusters. |
Interesting examples of faint cwarls are Boottes LL AL. = -3.1] 2007).. Coma AL. = -3.7] (Belokuroyet and Willman L ML. = -2.5] (Martinetal.2007). | Interesting examples of faint dwarfs are Boöttes II $_{v}$ = -3.1] \citep{ wjw07}, Coma $_{v}$ = -3.7] \citep{bel07} and Willman 1 $_{v}$ = -2.5] \citep{mar07}. |
. All of these galaxies are lainter than the overwhelming majority of globular clusters. | All of these galaxies are fainter than the overwhelming majority of globular clusters. |
The size of a globular clusters is best. described. by its half-light racius Ry. because this parameter remains almost invariant over ~ LO cluster relaxation times. | The size of a globular clusters is best described by its half-light radius $_{h}$, because this parameter remains almost invariant over $\sim$ 10 cluster relaxation times. |
The vast majority of elobular clusters have Ly« 10 pe. | The vast majority of globular clusters have $_{h} <$ 10 pc. |
However. some well-established globular clusters have quite large radii. | However, some well-established globular clusters have quite large radii. |
Examples are NGC 2419 (1t, = 18 pe) and Palomar 14 (It, = 25 pc). | Examples are NGC 2419 $_{h}$ = 18 pc) and Palomar 14 $_{h}$ = 25 pc). |
On the other hand Willman 1 (Martin 2007).. which has a hall-light radius of ~ 20 pe. has generally been regarded as a galaxy. | On the other hand Willman 1 \citep{mar07}, which has a half-light radius of $\sim$ 20 pc, has generally been regarded as a galaxy. |
Fhis overlap in size and luminosity raises deep questions about the nature of the distinction. between these two classes of objects. | This overlap in size and luminosity raises deep questions about the nature of the distinction between these two classes of objects. |
lt shoule be noted that the distinction. between galactic nuclei ancl globular clusters is also somewhat artificial. | It should be noted that the distinction between galactic nuclei and globular clusters is also somewhat artificial. |
Over time tidal forces will. lor cxample. strip away much of the stellar population of the Sagittarius dwarf spheroidal. leaving behind only its nucleus the luminous globular cluster AI54. | Over time tidal forces will, for example, strip away much of the stellar population of the Sagittarius dwarf spheroidal, leaving behind only its nucleus the luminous globular cluster M54. |
vandenBergh&Alackey(2004) ancl Alackey&vandenBereh(2005) have suggested: that he position of an object in a plot of M, versus log Ry, might be used to discriminate between elobular clusters and the stripped cores of dwarl galaxies. such as w Centauri. | \citet{vdbm04} and \citet{macvan05} have suggested that the position of an object in a plot of $_{v}$ versus log $_{h}$ might be used to discriminate between globular clusters and the stripped cores of dwarf galaxies, such as $\omega$ Centauri. |
However. this proposal now seems less attractive than it once did. | However, this proposal now seems less attractive than it once did. |
In particular vancenBergh(2007) has recently. noted. that the brightest red objects in the halo of the elliptical galaxy NGC 5128 appear to form a continuum in the M, versus log 1t,.(2007). | In particular \citet{vdb07} has recently noted that the brightest red objects in the halo of the elliptical galaxy NGC 5128 appear to form a continuum in the $_{v}$ versus log $_{h}$,. |
Furthermore NGC 2419. which on the basis of its position in the M, versus log 1t, plane. had. been classified: as a μαripped galaxy core. turns out to have a small metallicity lispersion and lacks a significante population of extra-tidal μαars (Bellazzini2007.Ripepictal.2007). | Furthermore NGC 2419, which on the basis of its position in the $_{v}$ versus log $_{h}$ plane, had been classified as a stripped galaxy core, turns out to have a small metallicity dispersion and lacks a significant population of extra-tidal stars \citep{bellrip07}. |
. Both of these factors militate against the hypothesis that NCC 2419 is wlually a stripped galaxy core. | Both of these factors militate against the hypothesis that NGC 2419 is actually a stripped galaxy core. |
On the other hand. the stripped. core suspect B514 in the Andromeda galaxy docs seer to be embedded in a very low surface density. dwarl μαoheroidal (Fredericiet.al.2007). | On the other hand the stripped core suspect B514 in the Andromeda galaxy does seem to be embedded in a very low surface density dwarf spheroidal \citep{fre07}. |
. La sumniary it appears wt the location of a galaxy in the Aly versus log Rh plot is not always a reliable way of separating globular clusters from the stripped cores of dwarl spheroidal galaxies. | In summary it appears that the location of a galaxy in the Mv versus log Rh plot is not always a reliable way of separating globular clusters from the stripped cores of dwarf spheroidal galaxies. |
The pioneering investigations of Babcock(1939). and Oort(1940). first demonstrated that dark matter provides a significant contribution to the masses of individual galaxies. | The pioneering investigations of \citet{bab39} and \citet{oor40}
first demonstrated that dark matter provides a significant contribution to the masses of individual galaxies. |
More recently. Mateo(1998). showed that such dark matter becomes more and more dominant as one proceeds to study ever dimmer galaxies. | More recently \citet{mat98} showed that such dark matter becomes more and more dominant as one proceeds to study ever dimmer galaxies. |
Ht has therefore become customary to regard the presence of dark matter as the touchstone that allows one to unambiguously distinguish. between galaxies and star clusters. | It has therefore become customary to regard the presence of dark matter as the touchstone that allows one to unambiguously distinguish between galaxies and star clusters. |
Furthermore. it is often dillicult and time consuming to obtain velocity dispersions of [aint stars in | Furthermore, it is often difficult and time consuming to obtain velocity dispersions of faint stars in |
scale for cosmological objects for which uo redshifts are availabe. | scale for cosmological objects for which no redshifts are available. |
lu his paper. we will deimoustrate a correlation between t1ο Euclidean valte of and the spectral hardness of GRBs. | In this paper, we will demonstrate a correlation between the Euclidean value of and the spectral hardness of GRBs. |
We interpret this correlationiin terius of a luiulnosity-harduess correlation but are iuitally unable to show the correlatiou exdlicitly. | We interpret this correlation in terms of a luminosity-hardness correlation but are initally unable to show the correlation explicitly. |
We use he - harduess correlation to derive the CRB luminosity [unction. as well as predicted counts as a function of flux aud redshift. | We use the - hardness correlation to derive the GRB luminosity function, as well as predicted counts as a function of flux and redshift. |
Finally. we show tle results of a siiiulation producir& the Correlation. | Finally, we show the results of a simulation producing the luminosity-hardness correlation. |
The paper is organized as follows. | The paper is organized as follows. |
In See. | In Sec. |
2. we presen the observed. co‘relation betweei spectal hardness and. | 2, we present the observed correlation between spectral hardness and. |
. The methodology. used o derive the luminosity Ancetion is cliscussecl in Sec. | The methodology used to derive the luminosity function is discussed in Sec. |
3. | 3. |
The resulting luminosity finetious are show‘in Sec. | The resulting luminosity functions are shown in Sec. |
|. togetler with predicte distrioutious of flux aud redshift. | 4, together with predicted distributions of flux and redshift. |
We show the results of a siuulation produciug the correlation in Sec. | We show the results of a simulation producing the luminosity-hardness correlation in Sec. |
5. lollowed by the discussion in Sec. | 5, followed by the discussion in Sec. |
6. | 6. |
Throughout this paper. we wil be usiug a [at cosmological model with Hy=65 kins + 1Oa; =0.3. and O4=0.7 1999). | Throughout this paper, we will be using a flat cosmological model with $H_0 = 65~$ km $^{-1}$ $^{-1}$, $\Omega_M = 0.3$, and $\Omega_{\Lambda} = 0.7$ \citep{bah99}. |
. In this paper we 1se a large homogeneous sample. the BD2 sample. derived. (roii. BATSE DISCLA data consisti1 of the continuous data stream from the eight BATSE LAD cetectors in four enerey channels «dL a timescale of 1021 us (Fislunanetal.1989). | In this paper we use a large homogeneous sample, the BD2 sample, derived from BATSE DISCLA data consisting of the continuous data stream from the eight BATSE LAD detectors in four energy channels on a timescale of 1024 ms \citep{fis89}. |
. The sample was derived uxing a software trigger alegorithin that interpolated the backgrouud between given times before and alter the ouset of the )ust aud required an excess of at least 5o over background in at least two detectors iu the energy range DO—3OO keV (Schmidt1999a). | The sample was derived using a software trigger algorithm that interpolated the background between given times before and after the onset of the burst and required an excess of at least $\sigma$ over background in at least two detectors in the energy range $50-300$ keV \citep{sch99a}. |
. The first version (he BDI sample) was described in οτι](t(1999a). | The first version (the BD1 sample) was described in \citet{sch99a}. |
. A revision discussed in Sec. | A revision discussed in Sec. |
2 of Schinidt(1999b) produced the BD2 sample. | 2 of \citet{sch99b}
produced the BD2 sample. |
The BD2 sample covers a period of 5.9 v from TJD 8365—10528. | The BD2 sample covers a period of 5.9 y from TJD $8365-10528$. |
It contains 1391 GRBs. olf which 1013 are also listed in the BATSE catalog. | It contains 1391 GRBs, of which 1013 are also listed in the BATSE catalog. |
The median photon flux limit of the BD? sample over the energy range 50—300 keV is 0.31 ρα7s +. | The median photon flux limit of the BD2 sample over the energy range $50-300$ keV is 0.31 ph $^{-2}$ $^{-1}$. |
The average Euclidean is 0.33640.008. | The average Euclidean is $0.336\pm0.008$. |
The sample of 1391 GRBs ellectively represents 2.00:) v of Full sky coverage. | The sample of 1391 GRBs effectively represents 2.003 y of full sky coverage. |
I Üistudying the correlation of spectral harduess of GRBs with other properties. we Lave to choose a relevant part of the light curve since the spectral hareuess of GRBs generally varies while thebust is going on. | In studying the correlation of spectral hardness of GRBs with other properties, we have to choose a relevant part of the light curve since the spectral hardness of GRBs generally varies while the burst is going on. |
The Euclidean values of in the BD2 sample have been derived from sinulations in which the CRB is moved out until the detectioi algorithin fails to trigger ).. | The Euclidean values of in the BD2 sample have been derived from simulations in which the GRB is moved out until the detection algorithm fails to trigger \citep{sch99a}. |
Iu most cases. the final detectiou is ou the peak of tie GRB time profile. | In most cases, the final detection is on the peak of the GRB time profile. |
Therefore. we use the 1021 tus interval containiug the peak flux to derive tle harduess ratio HR32 as the ratio of the burst couuts in BATSE channels 3 (100—300 keV) anc 2 (50—100 keV) for the brightest | Therefore, we use the 1024 ms interval containing the peak flux to derive the hardness ratio HR32 as the ratio of the burst counts in BATSE channels 3 $100-300$ keV) and 2 $50-100$ keV) for the brightest |
Fie. | Fig. |
3 presents our best estimate of the absorption-corrected. N-ray. [Inx of this object with time. based on these observations. (he papers of Dahlilem.Heckman&Fabbiano(1995).. Fabbiano.Heckman&Keel(1990).. Dahlemetal.(1996) Yaqoobetal.(1995)... and the count rate ina 2 ks ACIS-S3 observation taken on 1999 November 3. as part ola GTO program (Plak 2000. private communication). | \ref{fig:historical} presents our best estimate of the absorption-corrected X-ray flux of this object with time, based on these observations, the papers of \citet{dahlem95}, , \citet{fhk90}, \citet{dahlem96} \citet{yaqoob95}, and the count rate in a 2 ks ACIS-S3 observation taken on 1999 November 3, as part of a GTO program (Ptak 2000, private communication). |
We used the power law model that best-lits the Chandra spectrum to predict.Einstein.ROSAT (PSPC HRI) and count rates. | We used the power law model that best-fits the Chandra spectrum to predict, (PSPC HRI) and count rates. |
The ratio of the observed to predicted count rate. lor any particular observation. multiplied by the absorption corrected X-ray flux. (0.3 8.0 keV energv band) from this Chandra observation is the estimated. X-ray. flix plotted in Fig. 3.. For theEinstein. | The ratio of the observed to predicted count rate, for any particular observation, multiplied by the absorption corrected X-ray flux (0.3 – 8.0 keV energy band) from this Chandra observation is the estimated X-ray flux plotted in Fig. \ref{fig:historical}. |
PSPC and observations we have ealeulated a correction to obtain Che flux due to the bright source alone (given the significant Πας from other point sources and diffuse emission. see 4)). | For the, PSPC and observations we have calculated a correction to obtain the flux due to the bright source alone (given the significant flux from other point sources and diffuse emission, see \ref{sec:results:spectral}) ). |
We estimate that approximately of the count rate. of theROSAT PSPC count rate and of the count rate was due to diffuse X-ray emission and point sources other than the INO. | We estimate that approximately of the count rate, of the PSPC count rate and of the count rate was due to diffuse X-ray emission and point sources other than the IXO. |
The relative fluxes of the INO during the Einstein ROSAT and ASCA observations cliffer somewhat from the work of Dahlem (1995) and Dahlem.Weaver&Heckman(1993).. due to our correction for emission from other point sources and diffuse gas. and due to the strong dependence of absorption-corrected fIux on the assumed spectral shape. | The relative fluxes of the IXO during the Einstein ROSAT and ASCA observations differ somewhat from the work of Dahlem (1995) and \citet{dwh98}, due to our correction for emission from other point sources and diffuse gas, and due to the strong dependence of absorption-corrected flux on the assumed spectral shape. |
The IXO was substantially brighter in the early 1990s than it is now. with an intrinsic N-rav luminosity (0.3. 8.0 keV energy. band) of ~5x10Moros4. | The IXO was substantially brighter in the early 1990s than it is now, with an intrinsic X-ray luminosity (0.3 – 8.0 keV energy band) of $\sim 5 \times 10^{40} \ergps$. |
This makes it one of the most luminous ΕΝΟΣ. only slightly less luminous than the source in M82 (xaaret 2001: Makishima 2000). | This makes it one of the most luminous IXOs, only slightly less luminous than the source in M82 (Kaaret 2001; Makishima 2000). |
A drop in flix by a factor Z27 occurred between late 1991 Dec. and 1994 Mav (Dahlem 1995). | A drop in flux by a factor $\gtrsim 27$ occurred between late 1991 Dec. and 1994 May (Dahlem 1995). |
Our estimate of a low X-ray flux from the INO in the 1993 December observation is consistent with this [nding. given the non-detection of the source bx theROSAT IRI only 5 months later. | Our estimate of a low X-ray flux from the IXO in the 1993 December observation is consistent with this fading, given the non-detection of the source by the HRI only 5 months later. |
One should bear in mind that Fig. | One should bear in mind that Fig. |
3. is based on (he assumption that spectral shape ol the source has remained the same. only changing in absolute luminosity. | \ref{fig:historical} is based on the assumption that spectral shape of the source has remained the same, only changing in absolute luminosity. |
Given the large variations in derived luminosity. it is quile possible that the spectral shape has changed. due to inirimsic changes in the source itself or possibly even variations in foreground. absorption (as discussed bvDahlem.Heckman&Fabhiano (1995))). | Given the large variations in derived luminosity, it is quite possible that the spectral shape has changed, due to intrinsic changes in the source itself or possibly even variations in foreground absorption (as discussed by\citet{dahlem95}) ). |
The far-infrared (FIR) emission from galaxies is dominated by the thermal emission from dust grains (e.g.Harper&Low1973).. and is known to be a good indicator of star formation rate (e.g.Ken-nicutt1998:Inoue.Hirashita.&Kamaya 2000). | The far-infrared (FIR) emission from galaxies is dominated by the thermal emission from dust grains \citep[e.g.][]{harper73}, and is known to be a good indicator of star formation rate \citep[e.g.][]{kennicutt98,inoue00}. |
. Since the major heating source of dust grains is usually stellar radiation. the FIR colour. which reflects the dust temperature. is one of the most direct indicators of the interstellar radiation field (ISRF). | Since the major heating source of dust grains is usually stellar radiation, the FIR colour, which reflects the dust temperature, is one of the most direct indicators of the interstellar radiation field (ISRF). |
Moreover. the FIR emission is usually optically thin in the ISM. which means that we can trace all the radiation over the entire column. | Moreover, the FIR emission is usually optically thin in the ISM, which means that we can trace all the radiation over the entire column. |
(e.g.Walterbos&Greenawalt1996). 881 is particularly suitable for our purpose because (1) it is a nearby big galaxy with a face-on geometry: (1l) it can be compared with another well-known spiral galaxy. the Milky Way: and (iil) it is within a field of view of the sean observation ofAKAR/.. which is useful to quantify the global properties and to subtract the Which wavelengths in FIR are. necessary? | \citep[e.g.][]{walterbos96} 81 is particularly suitable for our purpose because (i) it is a nearby big galaxy with a face-on geometry; (ii) it can be compared with another well-known spiral galaxy, the Milky Way; and (iii) it is within a field of view of the scan observation of, which is useful to quantify the global properties and to subtract the Which wavelengths in FIR are necessary? |
At long wavelengths ἐν90 pum). the dust emission is mainly contributed by relatively large grains in radiative equilibrium with the stellar radiation field. | At long wavelengths $\ga 90~\micron$ ), the dust emission is mainly contributed by relatively large grains in radiative equilibrium with the stellar radiation field. |
However. an emission excess in shorter wavelengths ( 654m) is usually observed. which comes from very small dust grains suffering from stochastic heating. that is. heated transiently to a high temperature (Aannestad&Kenyon 1985). | However, an emission excess in shorter wavelengths $\la$ $\micron$ ) is usually observed, which comes from very small dust grains suffering from stochastic heating, that is, heated transiently to a high temperature \citep{aannestad79,draine85}. |
In order to estimate the large | In order to estimate the large |
EE μαμα eee αι —[000 παm M. | We plot the dispersion relation for both the long- and short-wavelength approximations at a value of $p_{\rm c}/k_{B}=1000$ K $^{-3}$ in Figure \ref{f9}. |
μα EE EE μαι πα πα παπα μα M IM IJI | The dotted and dashed lines are, respectively, the dispersion relations in the long- and short-wavelength approximations. |
ur | The solid line is a schematic dispersion relation expected to be obtained from an analysis without approximations. |
III παMEMNEIL I| πα παπα απ μα M NE É| —L NICE | From this figure, we expect the most unstable scale to be about twice the stabilized scale. |
CITNMUTIETDTMEIIEIPRIEMIITEENIUM | For such a scale, the growth timescale of the instability is comparable to the cooling timescale of the WNM $\sim 1$ Myr). |
AT IDA | The growth timescale can be shorter in a dynamical environment because of the effect of finite spatial extent as studied in 2.2. |
IT qo el | The evaporation flow speed becomes as large as $1$ km $^{-1}$. |
0)0, a T Lee -"- II | For this flow speed, from equation \ref{GR}) ), the growth timescale can be as small as the cooling time-scale of the CNM $\sim 0.1-0.01$ Myr). |
IM EENNM In Zl IIN |D X—-LOLLLÉLL. | Therefore, the instability is expected to grow within a dynamical timescale of the ISM $\sim 1-10$ Myr) and may drive a turbulent two-phase medium in the nonlinear stage, as in the case of a combustion front. |
LIL I I EMEIESM | Koyama Inutsuka (2002) and Audit Hennebelle (2005) have reported on the two-phase turbulence in numerical simulations of the thermally bistable medium, for which the velocity dispersion is on the order of the observed velocity dispersions in the diffuse ISM. |
IIII Im SEE SSR I MI α α α α Ἡ ππυπα μπα αἱ 1. BIBNMM μπα | They propose that turbulence in the ISM can be driven in a thermally unstable gas constantly supplied by shock propagation (Koyama Inutsuka 2002) or converging flows (Audit Hennebelle 2005; see also Heitsch et al. |
et al. | 2005; V\'azzquez-Semadeni et al. |
Subsets and Splits
No saved queries yet
Save your SQL queries to embed, download, and access them later. Queries will appear here once saved.