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where © is the solid angle subtended by the reprocessor from the Xray source (Done Navakshin 2001). | where $\Omega$ is the solid angle subtended by the reprocessor from the X–ray source (Done Nayakshin 2001). |
The amplitude of the iron. spectra eatures is further reduced by the Comptonization. | The amplitude of the iron spectral features is further reduced by the Comptonization. |
Two most important parameters determining the hickness of the hot laver are: the strength. of irradiation compared to internal disc emission. AN/fp. and the ratio of eravity at the base of the laver to Xray radiation pressure (RCOG. NIXIN). | Two most important parameters determining the thickness of the hot layer are: the strength of irradiation compared to internal disc emission, $\FX/\FD$, and the ratio of gravity at the base of the layer to X–ray radiation pressure (RC96, NKK). |
Following NlxIx we parameterize the latter nu where where f,=píngc1.9 [or cosmic abundance. | Following NKK we parameterize the latter by where where $\mu_{\rm m} = \rho/\nH \approx 1.9$ for cosmic abundance. |
Obviously then. the structure of the hot laver has to be solved. simultaneously with the structure of the cold. disc. since ct depends on the dise thickness. df. | Obviously then, the structure of the hot layer has to be solved simultaneously with the structure of the cold disc, since $A$ depends on the disc thickness, $\Hd$. |
Nit used vertically averaged solutions of Shakura Sunvacy (1973: hereafter. SS) for the disc in their computations of μμ. | NKK used vertically averaged solutions of Shakura Sunyaev (1973; hereafter SS) for the disc in their computations of $\tauh$. |
They assumed that total pressure. Poa.|Laat. should be continuous across the boundary between the two lavers. and obtained both 2... and Lea individually continuous across the boundary. | They assumed that total pressure, $\Pgas+\Prad$, should be continuous across the boundary between the two layers, and obtained both $\Pgas$ and $\Prad$ individually continuous across the boundary. |
In this paper we explicitly. solve equations of vertical structure of the cold. disc. | In this paper we explicitly solve equations of vertical structure of the cold disc. |
We do this for the a£, SS disc. aa disc with energv generation. proportional to total pressure. | We do this for the $\alpha\Ptot$ SS disc, a disc with energy generation proportional to total pressure. |
We define. the input parameters of the computations in such a wav as to make the connection between accretion discs models ancl observable quantities obvious. | We define the input parameters of the computations in such a way as to make the connection between accretion discs models and observable quantities obvious. |
First. we ceseribe calculations at a given. radius and then discuss radial dependence of 75,4. | First, we describe calculations at a given radius and then discuss radial dependence of $\tauh$. |
The structure of the LL at a given radius is computed. by combining the method of RCOG and ROO with computations of the vertical structure of X-ray illuminated. dises by Roaansska οἱ ((1999). | The structure of the HL at a given radius is computed by combining the method of RC96 and R99 with computations of the vertical structure of X-ray illuminated discs by R\'{o}\\.{z}aa\'{n}sska et (1999). |
For the LIL we follow. closely the method of C06 and ROO. with the important simplification of neglecting thermal conduction. | For the HL we follow closely the method of RC96 and R99, with the important simplification of neglecting thermal conduction. |
“Phe reason is) purely practical. as elforts to fully combine proper photo-ionization computations of the hot laver with vertical disc structure have only just begun (Dumont. Abrassart Collin 2000: ltó.zARSSEkA et alb. | The reason is purely practical, as efforts to fully combine proper photo-ionization computations of the hot layer with vertical disc structure have only just begun (Dumont, Abrassart Collin 2000; R\'{o}\\.{z}aa\'{n}sska et al., |
in preparation). | in preparation). |
This leaves us with an important. problem of selecting proper solution in the zone of instabilitv. where the svstem of equations has more than one solution for temperature. | This leaves us with an important problem of selecting proper solution in the zone of instability, where the system of equations has more than one solution for temperature. |
We adopt a simple prescription and select the highest value of Z as the solution. | We adopt a simple prescription and select the highest value of $T$ as the solution. |
Comparing solutions with and without conduction in RCOG and 1099 one sees that our procedure may overestimate somewhat the total thickness of the hot and transition lavers tthe depth of the point where 7= Zip). | Comparing solutions with and without conduction in RC96 and R99 one sees that our procedure may overestimate somewhat the total thickness of the hot and transition layers the depth of the point where $T=\Teff$ ). |
However. the thickness of the hot laver. alone is not. alfected by our neglecting the thermal conduction. | However, the thickness of the hot layer alone is not affected by our neglecting the thermal conduction. |
The spectrum. of illuminating radiation is assumed to be a power law with a eutoll. fexbtte7/7, parameterized by the photon spectral index E and cutoll energy. Leo. | The spectrum of illuminating radiation is assumed to be a power law with a cutoff, $F_E \propto E^{-\Gamma+1} {\rm e}^{-E/E_{\rm c}}$, parameterized by the photon spectral index $\Gamma$ and cutoff energy, $E_{\rm c}$. |
Equations of vertical structure of the cold cise are the same as in LGALaadsska et (1999). | Equations of vertical structure of the cold disc are the same as in R\'{o}\\.{z}aa\'{n}sska et (1999). |
We assume that a certain fraction of eravitational energy. £. is dissipated within the dise (but the disc transports all the angular momentum. see WWitt. Czerny Zvvcki 1997). | We assume that a certain fraction of gravitational energy, $\xi$, is dissipated within the disc (but the disc transports all the angular momentum, see Witt, Czerny Żyycki 1997). |
The remaining fraction. | £. is dissipated in an active corona and converted to hard X-ray radiation. of which a fraction 0.57) illuminates the disk. i.e. where is the gravitational energy. dissipation per unit area of the disc surface. | The remaining fraction, $1-\xi$ , is dissipated in an active corona and converted to hard X-ray radiation, of which a fraction $0.5 \eta$ illuminates the disk, i.e. where is the gravitational energy dissipation per unit area of the disc surface. |
Value of y=1 would thus correspond to a | Value of $\eta=1$ would thus correspond to a |
Exomoons. the satellites of extrasolar planets have been often featured in fiction as habitable locations. | Exomoons, the satellites of extrasolar planets, have been often featured in fiction as habitable locations. |
There is no deficit of known geiaut planets: Exoplanuct.ore (Wrightetal2010)0 lists approximately 10 eiaut exopluiets of total) within of the equilibriun temperature of Earth. as are 30 (3%)) of the planet candidates released in February 2011 (Doruclietal. 2011). | There is no deficit of known giant planets; Exoplanet.org \citep{ExoplanetOrg} lists approximately 40 giant exoplanets of total) within of the equilibrium temperature of Earth, as are 30 ) of the planet candidates released in February 2011 \citep{Borucki2011}. |
. Though these observations are preliminary. they do show that habitable-zoue eiaut planets not only exist. but are conuuon. | Though these observations are preliminary, they do show that habitable-zone giant planets not only exist, but are common. |
Once a egiaut planet is known to be in a habitable zone. variations du its orbit. such as Trausit Tiuiuug Variation (TTV:Sartorctti and Transit Duration Variation (ΤΟΝKippiug 2009).. photometry (Szabóetal.2006).. or gravitational imuücroleusiug (Liebig&4Wanbseganss 2010).. allow the indirect detection of satellites. | Once a giant planet is known to be in a habitable zone, variations in its orbit, such as Transit Timing Variation \citep[TTV;][]{Sartoretti1999} and Transit Duration Variation \citep[TDV;][]{Kipping2009}, , photometry \citep{Szabo2006}, or gravitational microlensing \citep{Liebig2010}, allow the indirect detection of satellites. |
Thus. if potentially habitable cxomoons exist around trausitiug eiut planets. they may be detected at the same (or even ereater) rate as solitary habitable terrestrial planets; | Thus, if potentially habitable exomoons exist around transiting giant planets, they may be detected at the same (or even greater) rate as solitary habitable terrestrial planets. |
As vet. uo exomoons have been detected. but the wealth of ransit data from the Nepler mission should beein to fill his gap. | As yet, no exomoons have been detected, but the wealth of transit data from the Kepler mission should begin to fill this gap. |
Despite the existence of giant plaucts dn stellar wabitable zones. it ds far from certain how they arrived there. | Despite the existence of giant planets in stellar habitable zones, it is far from certain how they arrived there. |
Current giaut-plauet formation models asstune that they are created at distances bevoud 1ο stability. poiut of ice (e.g.LissauerL987:Boss 1997).. which miplies conditions not suitable to surface wabitability. | Current giant-planet formation models assume that they are created at distances beyond the stability point of ice \citep[e.g.][]{Lissauer1987,Boss1997}, which implies conditions not suitable to surface habitability. |
Disk uueration can bring giaut planets close to the star (Ward1997).. but ecnerally has a stopping point far too close to the star to be habitable (thus producing Πο Jupiters”). | Disk migration can bring giant planets close to the star \citep{Ward1997}, but generally has a stopping point far too close to the star to be habitable (thus producing "Hot Jupiters"). |
The host planets of potentially habitable exomoous therefore likely arrived at their final orbit through late-stage migration. driveu either by planctesimals (Ixirslietal.2009) or other eiaut ylancets (Weidensclilling&\arzazi1996). | The host planets of potentially habitable exomoons therefore likely arrived at their final orbit through late-stage migration, driven either by planetesimals \citep{Kirsh2009} or other giant planets \citep{Weidenschilling1996}. |
. Tn the process of nügrating. the satellite svstenis of these giaut planet may have close encouuters with errestrial plauets or planetesimals. causing them to be disrupted or replaced. | In the process of migrating, the satellite systems of these giant planet may have close encounters with terrestrial planets or planetesimals, causing them to be disrupted or replaced. |
If either the Jovian or θαΕπστὰ systems were transported to L AU around a solar nass-star. both Callisto aud Titan would be at of their planets Till radii. thus implying that all the uajor satellites of the two planets would be on stable orbits. | If either the Jovian or Saturnian systems were transported to 1 AU around a solar mass-star, both Callisto and Titan would be at of their planet's Hill radii, thus implying that all the major satellites of the two planets would be on stable orbits. |
However. a close eucounuter could either excite heir orbits to ligh eccentricity (thus requiring tidal recizcularization). or could result iu the capture of a uuch larger terrestrial satellite. | However, a close encounter could either excite their orbits to high eccentricity (thus requiring tidal recircularization), or could result in the capture of a much larger terrestrial satellite. |
Neptune appears have to experienced this process during its mieration through the xoto-Ixuiper Belt. loosing any. original major satellites. while eaining Triton iun au inclined. retrograde orbit. | Neptune appears have to experienced this process during its migration through the proto-Kuiper Belt, loosing any original major satellites, while gaining Triton in an inclined, retrograde orbit. |
This was possibly due to a 1iomientuu-excehlauge reaction hat ejected the binazy companion of Triton (Agnor&Uaimiltou 2006).. though other scenarios are possible (at reduced: probabilitv). | This was possibly due to a momentum-exchange reaction that ejected the binary companion of Triton \citep{Agnor2006}, though other scenarios are possible (at reduced probability). |
Αν capture process. though. will tend to produce very. loosely-bouud initial orbits. with ouly a small delta-v to escape velocity at periapsc. | Any capture process, though, will tend to produce very loosely-bound initial orbits, with only a small delta-v to escape velocity at periapse. |
Therefore. some method ust be used to determine the one-term evolution and stability (or lack thereof) for hese orbits. | Therefore, some method must be used to determine the long-term evolution and stability (or lack thereof) for these orbits. |
Tere we use a full οΤΕ (Isozai Cyele aud Tidal Friction) model to find the survival probability for a ranee of plysical conditions aud the detectability of he resulting svsten. | Here we use a full KCTF (Kozai Cycle and Tidal Friction) model to find the survival probability for a range of physical conditions and the detectability of the resulting system. |
As shown bv Donuison(2010) and Sato&Asada (2010).. there are a range of stable orbits for Earthauass auets around eiaut plancts. | As shown by \citet{Donnison2010} and \citet{Sato2010}, there are a range of stable orbits for Earth-mass planets around giant planets. |
However. both of those uodels test ouly the stability of the orbit. rather than any evolution due to tidal effects. | However, both of those models test only the stability of the orbit, rather than any evolution due to tidal effects. |
On inclined exomoon orbits. the effects of stellar torques on the orbit can. hrough initiating I&ozai cvcles. dramatically accelerate he rate of tidal decay. orbit circularization. and spiu- svnchronization. | On inclined exomoon orbits, the effects of stellar torques on the orbit can, through initiating Kozai cycles, dramatically accelerate the rate of tidal decay, orbit circularization, and spin-orbit synchronization. |
As we will show. this process cau allow evenvery loose. inclined capture orbits to stabilize | As we will show, this process can allow evenvery loose, inclined capture orbits to stabilize |
this. in fact. is seen clearly in most cases in the colour index images. ancl also often in the broad-band images. especially theST NICMOS images. | this, in fact, is seen clearly in most cases in the colour index images, and also often in the broad-band images, especially the NICMOS images. |
Such dust lanes are expected to be continuations of the dust lanes in the bar. ancl this is wha is seen in our images (see Shlosman 1999 for a theoretica review). | Such dust lanes are expected to be continuations of the dust lanes in the bar, and this is what is seen in our images (see Shlosman 1999 for a theoretical review). |
ΝΑ. colour index images are not very sensitive to changes in stellar populations. and can outline clus structure clearly (see e.g. the ἐν map of MIOO in Ixnapen et al. | NIR colour index images are not very sensitive to changes in stellar populations, and can outline dust structure clearly (see e.g. the $I-K$ map of M100 in Knapen et al. |
1995a). | 1995a). |
. We see clear ancl abundant observationa evidence for dust lanes on several scales. but most clearly in the (Nits. | We see clear and abundant observational evidence for dust lanes on several scales, but most clearly in the CNRs. |
Dust lanes in the bars are not well visible in general in our NIHU imaging due to the lower signal to noise ratios achieved in the bar regions. | Dust lanes in the bars are not well visible in general in our NIR imaging due to the lower signal to noise ratios achieved in the bar regions. |
In. Paper IH. we present optical colour index maps which outline the dust lane structure in the bars and dises of our sample galaxies more clearly. and we study. the relationship between the shape of the dust lanes. the axial ratio. and the SE in the bar in more cletail. | In Paper II, we present optical colour index maps which outline the dust lane structure in the bars and discs of our sample galaxies more clearly, and we study the relationship between the shape of the dust lanes, the axial ratio, and the SF in the bar in more detail. |
Our colour maps have shown the presence of cireumnuclear rings in most of the galaxies in our sample. | Our colour maps have shown the presence of circumnuclear rings in most of the galaxies in our sample. |
These rings are generally redder than other regions in the galaxy. by about 0.05. 0.1 magnitude in 4A and 44A. | These rings are generally redder than other regions in the galaxy by about 0.05– 0.1 magnitude in $J-K$ and $H-K$. |
From the colour index maps or colour profiles alone it is not possible to make meaningful statements about quantities of extinguishing dust implied by the redcder colours. | From the colour index maps or colour profiles alone it is not possible to make meaningful statements about quantities of extinguishing dust implied by the redder colours. |
In fact. there are indications from both optical and. NLR imaging and spectroscopy that the red. colours in CNRs in galaxies like the ones studied here may be inlluenced by voung stars. e.g. red supergiants (Ixnapen et al. | In fact, there are indications from both optical and NIR imaging and spectroscopy that the red colours in CNRs in galaxies like the ones studied here may be influenced by young stars, e.g. red supergiants (Knapen et al. |
1995a.b: Ixnapen. 1996: Elmeercen et al. | 1995a,b; Knapen 1996; Elmegreen et al. |
1997: Rweler Ixnapen 1999). | 1997; Ryder Knapen 1999). |
In Paper LI. we will compare the precise location of the JA and dfA features with those of SE regions as seen in comission. and try to place quantitative limits on the origins of the red light in the ςΝις. | In Paper III, we will compare the precise location of the $J-K$ and $H-K$ features with those of SF regions as seen in emission, and try to place quantitative limits on the origins of the red light in the CNRs. |
Two of the host galaxies of Sy nuclei (NGC 3516 and NGC 3982). which have colours consistent with those nmieasured by Peletier et al. ( | Two of the host galaxies of Sy nuclei (NGC 3516 and NGC 3982), which have colours consistent with those measured by Peletier et al. ( |
1999). also show a peculiar shape in the colour profiles. starting from a very red. value anc steeply decreasing until the radius of the ring. remaining constant afterwards. | 1999), also show a peculiar shape in the colour profiles, starting from a very red value and steeply decreasing until the radius of the ring, remaining constant afterwards. |
Peletier ct al. ( | Peletier et al. ( |
1999) suggested. tha the red. LfLy (or J dv) colours in the cores of many ον galaxies could be due to a significant [fraction of therma radiation [rom hot dust heated by the Sy nucleus. | 1999) suggested that the red $H-K$ (or $J-K$ ) colours in the cores of many Sy galaxies could be due to a significant fraction of thermal radiation from hot dust heated by the Sy nucleus. |
However. the two other ACN hosts in our sample (NGC 4303. aux NGC 6951) do not show significantly red nuclei. | However, the two other AGN hosts in our sample (NGC 4303 and NGC 6951) do not show significantly red nuclei. |
dispersion iu the cluster aud a characteristic radius. | dispersion in the cluster and a characteristic radius. |
Figure TU shows a portion of the NIRSPEC echelle spectral of M8S2-F. compared to a series of template supereiaut spectra. | Figure \ref{spect} shows a portion of the NIRSPEC echelle spectrum of M82-F, compared to a series of template supergiant spectra. |
Features found in the spectra of supereiant stars are readilv ideuti&ed iu the cluster spectrum. although hey appear "washed οπ due to the stellar velocity dispersion. | Features found in the spectra of supergiant stars are readily identified in the cluster spectrum, although they appear “washed out” due to the stellar velocity dispersion. |
Especially prominent are the rovibrational CO vandheads and umucrous Fe and ΟΠ lines. | Especially prominent are the rovibrational CO bandheads and numerous Fe and OH lines. |
We have assembled an atlas of high-resolution NIRSPEC spectra of supergiant stars for use in cross-correlation analysis (?).. roni which we determine the dominant spectral type and inc-ofsight velocity dispersion. | We have assembled an atlas of high-resolution NIRSPEC spectra of supergiant stars for use in cross-correlation analysis \citep{mccrady03}, from which we determine the dominant spectral type and line-of-sight velocity dispersion. |
Dased upon the peak of he cross-correlation function. (Fieure ??)). the IT-baud spectruni of MS82-F ποτ closely matches spectral types in he range of NUMOL. The lne-ofsight velocity dispersion owed on cross-correlation with templates in this rauge iso,=13.5x02 kins |. | Based upon the peak of the cross-correlation function (Figure \ref{ccfplot}) ), the $H$ -band spectrum of M82-F most closely matches spectral types in the range of K4I–M0I. The line-of-sight velocity dispersion based on cross-correlation with templates in this range is $\sigma_r = 13.5 \pm 0.2$ km $^{-1}$. |
The value of e, decreases as a function of the simibuitv to the template spectra as measured bv the peak value of the cross-correlation function. | The value of $\sigma_r$ decreases as a function of the similarity to the template spectrum as measured by the peak value of the cross-correlation function. |
It is therefore possible that the stated value reflects some template uvisimatch bias due to our limited teiiplate spectra atlas. but this effect is snall (< 0.5 kau S1 |. | It is therefore possible that the stated value reflects some template mismatch bias due to our limited template spectra atlas, but this effect is small $<$ 0.5 km $^{-1}$ ). |
The simplest approach to determine the cluster mass is to measure the haltleht radius. assume that light traces mass. the cluster is spherical. and the velocity dispersion is isotropic. then apply the virial theorem. | The simplest approach to determine the cluster mass is to measure the half-light radius, assume that light traces mass, the cluster is spherical, and the velocity dispersion is isotropic, then apply the virial theorem. |
\[82-F preseuts a difficulty for this method: IST iuages (Figure ??)) clearly show that the cluster is elliptical in projection aud cannot. therefore. be spherical. | M82-F presents a difficulty for this method: HST images (Figure \ref{fitplot}) ) clearly show that the cluster is elliptical in projection and cannot, therefore, be spherical. |
To measure the radius. we fit the cluster using au elliptical version of the empirical ? model: where e and 6 are the characteristic leusths of tle minor and major axes. respectively. A4 is a scaling constant and Ao acts as a “tidal radius truncating the profile bevond a particular scale leneth. | To measure the radius, we fit the cluster using an elliptical version of the empirical \citet{king62} model: where $a$ and $b$ are the characteristic lengths of the minor and major axes, respectively, $k_1$ is a scaling constant and $k_2$ acts as a “tidal radius,” truncating the profile beyond a particular scale length. |
The fit therefore has four free parameters (Ay. Ko. à and b) to describe the liebt profile aud three more to describe the ceutroid location and position anele. | The fit therefore has four free parameters $k_1$, $k_2$, $a$ and $b$ ) to describe the light profile and three more to describe the centroid location and position angle. |
The ine mocel is convolved with a model PSF from Tiny Tiu (7)| and compared to the Πμαρσο | The King model is convolved with a model PSF from Tiny Tim \citep{krist95} and compared to the image. |
, Fit paraincters are determined by iterative search over parameter space. using a Levenbure-\larquardt least-squares fit. | Fit parameters are determined by iterative search over parameter space, using a Levenburg-Marquardt least-squares fit. |
Fieure ?? shows the fit aud residuals for the ACS/IIRC FalIW image of M82-F. The halfleht radius iu projection along the major axis. ppys is the semudniajor axis of the ellipse that eucloses half the ux in the fitted ing model. | Figure \ref{fitplot} shows the fit and residuals for the ACS/HRC F814W image of M82-F. The half-light radius in projection along the major axis, $r_{hp}$, is the semimajor axis of the ellipse that encloses half the flux in the fitted King model. |
We determine 7j, munerically by stuum4ne the fiux iu a series of ellipses with the axial ratio defined by «à aud b. | We determine $r_{hp}$ numerically by summing the flux in a series of ellipses with the axial ratio defined by $a$ and $b$. |
Moute Carlo simulations of clusters indicate that a. b aud fy are siguifcantlv covariant: however. the fitted. projected lalflight radius along a given axis is accurate to about 2 percent. | Monte Carlo simulations of clusters indicate that $a$, $b$ and $k_2$ are significantly covariant; however, the fitted projected half-light radius along a given axis is accurate to about 2 percent. |
The halfinass radius may be determined by assume that light traces niass aud deprojectiug by dividing rj, by 0.766 (7). | The half-mass radius may be determined by assuming that light traces mass and deprojecting by dividing $r_{hp}$ by 0.766 \citep{spitzer87}. |
. ? fit a spherical Nine model to MS2-F in the NICAIOS FLGOW nuage and found a projected halflelt radius of SO+1L mas. | \citet{mccrady03} fit a spherical King model to M82-F in the NICMOS F160W image and found a projected half-light radius of $89 \pm 11$ mas. |
Our elliptical Ning function fit to the same iuage found rj,=113X2 along the major axis aud rp=62+] along the minor axis. | Our elliptical King function fit to the same image found $r_{hp}=113 \pm 2$ along the major axis and $r_{hp}=62 \pm 1$ along the minor axis. |
At the adopted distance of M82. the projected half-light radi along the major and minor axes are thus 1.97£0.16 pe and 1140.1 pc. respectively, | At the adopted distance of M82, the projected half-light radii along the major and minor axes are thus $1.97 \pm
0.16$ pc and $1.1 \pm 0.1$ pc, respectively. |
To account for the ellipticity of the cluster. we assunie it is an oblate spheroid and compute the eravitational potential assuming that the cluster is homocoidal (7.Section2.3).. | To account for the ellipticity of the cluster, we assume it is an oblate spheroid and compute the gravitational potential assuming that the cluster is homoeoidal \citep[][Section 2.3]{binney87}. . |
The mmuerical constaut iu the virial mass formmla may be separated iuto factors dependent upon the central concentration aud the ellipticitv. | The numerical constant in the virial mass formula may be separated into factors dependent upon the central concentration and the ellipticity. |
The compact core plus extended euvelope structure of star clusters is simular to au η—5 polvtrope (7.p.13).. | The compact core plus extended envelope structure of star clusters is similar to an $n=5$ polytrope \citep[][p. 13]{spitzer87}. |
For anv —5 polvtropic oblate ellipsoid with au isotropic velocity dispersion. the virial mass ds: where the eccentricity is 6=Vl(Z/R)\2. and R and Z are the equatorial and polar radii of the oblate spheroid. respectively, | For an $n=5$ polytropic oblate ellipsoid with an isotropic velocity dispersion, the virial mass is: where the eccentricity is $e = \sqrt{1-(Z/R)^2}$, and $R$ and $Z$ are the equatorial and polar radii of the oblate spheroid, respectively. |
Figure ?? shows the fitted cluster lalf-light radii iu projection along the major and minor axes for the NICMOS aud ACS images. | Figure \ref{halfradii} shows the fitted cluster half-light radii in projection along the major and minor axes for the NICMOS and ACS images. |
The observed axial ratio. afb. veprescuts a lower bound ou the ratio Z/Ri: an oblate cluster will appear rouuder than its intrinsic shape uuless viewed directly aloug the equatorial plauc. | The observed axial ratio, $a/b$ , represents a lower bound on the ratio $Z/R$; an oblate cluster will appear rounder than its intrinsic shape unless viewed directly along the equatorial plane. |
Averaging over all possible iucliuatious. we find that the observed. axial ratio of 0.55 corresponds to a imost-likelv intrinsic axial ratio of 0.35 with e=(0.77x0.07. | Averaging over all possible inclinations, we find that the observed axial ratio of $0.55$ corresponds to a most-likely intrinsic axial ratio of $0.35$ with $e=0.77 \pm 0.07$. |
We represent the uucertaintv on the viewing inclination as the difference )etween the augle-averaged value aud the lower bouud. | We represent the uncertainty on the viewing inclination as the difference between the angle-averaged value and the lower bound. |
We believe the ellipticity of the cluster is intrinsic. rather hau the result of differcutial reddening. | We believe the ellipticity of the cluster is intrinsic, rather than the result of differential reddening. |
Extinction is very low at 2.2 gan. aud therefore the observed shape in the F222\0 nuages is verv likely intrinsic. | Extinction is very low at 2.2 $\mu$ m, and therefore the observed shape in the F222M images is very likely intrinsic. |
\oreover. he axial ratio is constant within the uncertainties across all wavebands frou 0.1 to 2.2 gan (Figure ??)). | Moreover, the axial ratio is constant within the uncertainties across all wavebands from 0.4 to 2.2 $\mu$ m (Figure \ref{halfradii}) ). |
If the ellipticitv were due to the distribution of dust arouud he cluster. the axial ratio should chiuge as a functiou of wavelength as the extinction ids expected to vary sienificantly between P aud K. | If the ellipticity were due to the distribution of dust around the cluster, the axial ratio should change as a function of wavelength as the extinction is expected to vary significantly between $B$ and $K$. |
The axial ratio does appear o increase sharply at the shortest IIST/ACS waveband (F250W). | The axial ratio does appear to increase sharply at the shortest HST/ACS waveband (F250W). |
We interpret this as a result of scatteriug of cluster Light by dust outside the cluster. | We interpret this as a result of scattering of cluster light by dust outside the cluster. |
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