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However. this correlation - if in fact correct - must also have significant scatter. as the most massive halo (#288) shows the lowest destruction rate. | However, this correlation - if in fact correct - must also have significant scatter, as the most massive halo 8) shows the lowest destruction rate. |
This might be linked to the environment and supply of new satellites. respectively. | This might be linked to the environment and supply of new satellites, respectively. |
‘Table 1 summarises the basic properties of our halos aud their environments. illustrating the variety of richness and accretion histories sampled. by our simulations. | Table \ref{HaloDetails} summarises the basic properties of our halos and their environments, illustrating the variety of richness and accretion histories sampled by our simulations. |
The mass spectra of the satellite galaxies (although not. presented) | The mass spectra of the satellite galaxies (although not presented) |
Wide field tagging is becoming imereasnglv conmuuon since new large format CCD cameras are. or soon will be available at several telescope (see c.g, ATEGACAAL Boulade 1998. Boulade et al. | Wide field imaging is becoming increasingly common since new large format CCD cameras are, or soon will be available at several telescope (see e.g. MEGACAM, Boulade 1998, Boulade et al. |
1998. WFI. Daade 1999. otc). | 1998, WFI, Baade 1999, etc). |
The possibility. to perform: wide field) imagine of the extragalactic sky allows a systematic search of meditu-hieh redshift ealaxw clusters iu two-dimensional photometric catalogs of galaxies. | The possibility to perform wide field imaging of the extragalactic sky allows a systematic search of medium-high redshift galaxy clusters in two-dimensional photometric catalogs of galaxies. |
These caudidate clusters are of cosinological interest and are primary targets for subsequent follow-up spectroscopical observations (sec. e.g.. Holden et al. | These candidate clusters are of cosmological interest and are primary targets for subsequent follow-up spectroscopical observations (see, e.g., Holden et al. |
1999. Riunella et al. | 1999, Ramella et al. |
2000). | 2000). |
Several automated aleorithnis have already been developed. for the detection of clusters within two dimensional galaxy catalogs. | Several automated algorithms have already been developed for the detection of clusters within two dimensional galaxy catalogs. |
The cassical techuiques used for this task are the "box count” (Lidiuau Peterson 1996) and the "natehed filter” algorithlii proposed bv Postinan ot al. ( | The classical techniques used for this task are the “box count” (Lidman Peterson 1996) and the “matched filter” algorithm proposed by Postman et al. ( |
1996. hereafter P96) auc its recent refinements (see Ixepuer et al. | 1996, hereafter P96) and its recent refinements (see Kepner et al., |
1999. Iiwvasald et al. | 1999, Kawasaki et al. |
1998. Lobo et al 1999). | 1998, Lobo et al 1999). |
The box-counting method. uses sliding windows (usualbiv snares} which are moved across the poiut distribution marking the positious where the couut rate in the central part of the window exceeds the value expected from the backerouud determined in the outermost regions of the window. | The box-counting method uses sliding windows (usually squares) which are moved across the point distribution marking the positions where the count rate in the central part of the window exceeds the value expected from the background determined in the outermost regions of the window. |
The ndn drawbacks of the method are the introduction of a binning to determine the local ckeround. which improves count statistics at the expense of spatial accuracy. and the depeudeuce on artificial paraneters like bin sizes aud positions or window size and geometry. | The main drawbacks of the method are the introduction of a binning to determine the local background, which improves count statistics at the expense of spatial accuracy, and the dependence on artificial parameters like bin sizes and positions or window size and geometry. |
The vmatched filter" is a maxiuumn-likelihood. (AIL) algoritlan which analyzes the galaxy distribution with the assunrptiou of some model profiles to fit the data (e.g. a density distribution profile aud a luninosity fiction). | The “matched filter” is a maximum-likelihood (ML) algorithm which analyzes the galaxy distribution with the assumption of some model profiles to fit the data (e.g. a density distribution profile and a luminosity function). |
This last technique has been used to build the Palomar Distaut Cluster Survey (PDCS. see P96) catalog and he EIS cluster catalog (Olsen ct al. | This last technique has been used to build the Palomar Distant Cluster Survey (PDCS, see P96) catalog and the EIS cluster catalog (Olsen et al. |
1999. Scodecgeio et al. | 1999, Scodeggio et al. |
1999). | 1999). |
These two catalogs are two of the largest xeseutlv available sets of distant clusters. with 79 aud 302 candidate clusters respectively. | These two catalogs are two of the largest presently available sets of distant clusters, with 79 and 302 candidate clusters respectively. |
However. the main drawback of the matched filter nethod is that it can muss clusters tha ave not sviuimetric or that differ siguificautly from the assumed profile. | However, the main drawback of the matched filter method is that it can miss clusters that are not symmetric or that differ significantly from the assumed profile. |
This can be a serious problem since we know that a large yaction of clusters have a pronounced ellipticity (see e.g. | This can be a serious problem since we know that a large fraction of clusters have a pronounced ellipticity (see e.g. |
The convergence criterion for the algorithm was taken to be when οἳ changed bv less than 10.1 at all radii. | The convergence criterion for the algorithm was taken to be when $\varrho^{*}$ changed by less than $10^{-4}$ at all radii. |
The computer program developed to implement this algorithm. written in C. uses a radial grid. H;. with points spaced logarithmically between an inner radius of 2=10 and an outer of 2=10°. | The computer program developed to implement this algorithm, written in C, uses a radial grid $R_i$, with points spaced logarithmically between an inner radius of $R = 10^{-5}$ and an outer of $R = 10^5$. |
The various properties of the svstem are described on this grid: 0;=ORs). e;=eG). and so on. | The various properties of the system are described on this grid; $\varrho_i = \varrho(R_i)$, $\psi_i = \psi(R_i)$, and so on. |
The DF is described on a grid of energv and angular momentum points. where the energv points have lor convenience the values of the potential on the radial grid. 2;= c. ancl the angular momentum is evaluated as v=j/j.. where j. is the circular angular momentum ancl. goes from zeroto one. | The DF is described on a grid of energy and angular momentum points, where the energy points have for convenience the values of the potential on the radial grid, $\varepsilon_i = \psi_i$ , and the angular momentum is evaluated as $x = j/j_c$, where $j_c$ is the circular angular momentum and $x$ goes from zeroto one. |
where Note that the elfect on the central intensity of the seeing convolved distribution is a function of the intrinsic ellipticity of the object. | where Note that the effect on the central intensity of the seeing convolved distribution is a function of the intrinsic ellipticity of the object. |
Basically. as € increases the spread of photons from the inner parts of the profile due to the seeing is more ellicient ancl consequentIy the central intensity decreases. | Basically, as $\epsilon$ increases the spread of photons from the inner parts of the profile due to the seeing is more efficient and consequently the central intensity decreases. |
In the absence of seeing. by construction. all isophotes of the profile have the same cllipticity. whereas the presence of secing tends to make them circular. | In the absence of seeing, by construction, all isophotes of the profile have the same ellipticity, whereas the presence of seeing tends to make them circular. |
Using the isophote condition. Z.(£.0)=£.(£.5/2) Ht is possible to derive an implicit equation that gives the variation of the ellipticity with the racial distance: For this problem. it is useful to introduce a,.«0? and i as functions of the Cartesian coordinates wry: and quom this implicit equation we can obtain gf.r and }Aherefore the ellipticity of the isophotes alfected. by sceing using εξ10gf. | Using the isophote condition, $I_{\rm c}(\xi,0)=I_{\rm c}(\xi,\pi/2)$ it is possible to derive an implicit equation that gives the variation of the ellipticity with the radial distance: For this problem, it is useful to introduce $a_\epsilon, w,
a_{\epsilon}^*$ and $w^*$ as functions of the Cartesian coordinates $x,y$: and From this implicit equation we can obtain $y/x$ and therefore the ellipticity of the isophotes affected by seeing using $\epsilon(x)=1-y/x$. |
The theory of atmospheric turbulence: predicts the PSE to be the Fourier transform of (by 4] (Fried. 1966: Woolf 1982).OS2). where EWLLALAl —= 292070060 iand b iss a scaling parameter. | The theory of atmospheric turbulence predicts the PSF to be the Fourier transform of $-(kb)^{5/3}$ ] (Fried 1966; Woolf 1982), where FWHM = $b$ and $b$ is a scaling parameter. |
In theread domain this PSE is written as: where Jy is the standard Bessel function. (Abramowitz Stegun 1964. p. 358). | In the domain this PSF is written as: where $J_0$ is the standard Bessel function (Abramowitz Stegun 1964, p. 358). |
For a given PWHAL we have evaluated he value of 3 that minimizes the dillerence between the xediction of the atmospheric turbulence theory and. the Moffat function by minimizing the x7 of thefit between both Shs. | For a given FWHM, we have evaluated the value of $\beta$ that minimizes the difference between the prediction of the atmospheric turbulence theory and the Moffat function by minimizing the $\chi^2$ of the fit between both PSFs. |
An optimum value of 3~4.765 was found. | An optimum value of $\beta\sim 4.765$ was found. |
In Fig. | In Fig. |
2 we have shown the cillerence between the PSE prediction rom turbulence theory and a Molfat. function for a value of j—4.765. | 2 we have shown the difference between the PSF prediction from turbulence theory and a Moffat function for a value of $\beta=4.765$. |
Ht can be seen that the agreement is quite good. | It can be seen that the agreement is quite good. |
A Mollat function could therefore be used to reliably mocdel the urbulence prediction. although. the PSL's usually measures in real images have bigger “wings”. or equivalently smaller values of 3. than those expected [rom the turbulence theory (c.g. Saelia et al. | A Moffat function could therefore be used to reliably model the turbulence prediction, although, the PSFs usually measured in real images have bigger “wings”, or equivalently smaller values of $\beta$, than those expected from the turbulence theory (e.g. Saglia et al. |
1903). | 1993). |
This is because the real seeing no only depends on atmospheric conditions but is also causec w imperfections in telescope optics. | This is because the real seeing not only depends on atmospheric conditions but is also caused by imperfections in telescope optics. |
The presence of these bigger “wines” in real images makes Molfat. functions a better choice to model the PSE han the turbulence theory prediction. | The presence of these bigger “wings” in real images makes Moffat functions a better choice to model the PSF than the turbulence theory prediction. |
As an example of his. current packages of data reduction in LRAL suggest a default value of;= 2.5. In order to span the range | As an example of this, current packages of data reduction in IRAF suggest a default value of $\beta=2.5$ In order to span the range |
"Therefore. referring to Leonardetal.2003.. we estimate the systematic uncertainty of the EPM-based distance as 60 that means £1 Mpe for SN 2002ap. | Therefore, referring to \cite{leo3}, we estimate the systematic uncertainty of the EPM-based distance as 60 that means $\pm 4$ Mpc for SN 2002ap. |
The total (raudoni plus systematic) error of the distance becomes £15 Mpc. or about TO%. | The total (random plus systematic) error of the distance becomes $\pm 4.5$ Mpc, or about 70. |
. The consisteney between the EPNE-distauce aud the other photometric distances of M71 may support the applicability of the EPAL for hwvperuovae. | The consistency between the EPM-distance and the other photometric distances of M74 may support the applicability of the EPM for hypernovae. |
However. it is nuportaut to note that Leonardetal.2005— found siuillu agreement between the EPAI-distance and he vbrightest red supereiants” (DRSC) distance (Sohu&Davidge. 1998)) for NGC 1637, | However, it is important to note that \cite{leo3} found similar agreement between the EPM-distance and the “brightest red supergiants” (BRSG) distance \cite{sohn2}) ) for NGC 1637. |
Among he distance measurement iethods discussed by Leonardetal.2003.. the BRSG technique is the only one hat gives cousisteut result with the EPAL | Among the distance measurement methods discussed by \cite{leo3}, the BRSG technique is the only one that gives consistent result with the EPM. |
Therefore. his agreement does nof. unfortunately, reduce the large systematic uucertaimtv of the EPM-distauce for MT£. | Therefore, this agreement does not, unfortunately, reduce the large systematic uncertainty of the EPM-distance for M74. |
We conclude that our first attempt to apply the EPALechnique to a hyperuova resulted in distance estimate for SN 2002ap that has remarkably good internal precision (good linearity. low scattering in Fie.10.) | We conclude that our first attempt to apply the EPM-technique to a hypernova resulted in distance estimate for SN 2002ap that has remarkably good internal precision (good linearity, low scattering in Fig.10.) |
Ou the other rand. the systematic uncertaity of the EPA\Lcistauce is quite hieh due to the problematic issue of the dilutiou actors as well as the complexity of hivperuova spectra at the carly phases. | On the other hand, the systematic uncertainty of the EPM-distance is quite high due to the problematic issue of the dilution factors as well as the complexity of hypernova spectra at the early phases. |
Note that the EPAI-distance is A6 AIpc lower than the one used iu previous papers (7.3 Mpc). thus. the inferred pliysical parameters of SN 2002ap nav need some revision. if we adopt this new distance. | Note that the EPM-distance is 0.6 Mpc lower than the one used in previous papers (7.3 Mpc), thus, the inferred physical parameters of SN 2002ap may need some revision, if we adopt this new distance. |
For example. the nickel mass may be somewhat lower than derived previously (οιο, Mazzalietal.2002:: Daudevetal. 2002:: Foleyetal. 2003)). | For example, the nickel mass may be somewhat lower than derived previously (e.g. \cite{mazz}; \cite{pande}; \cite{foley}) ). |
Some basic physical parameters of the SN explosious can be estimated by comparing the observations with model computations of the lieht curves and/or spectra. | Some basic physical parameters of the SN explosions can be estimated by comparing the observations with model computations of the light curves and/or spectra. |
This usually requires sophisticated calculations iucludiug lydvodvuaiics. radiative transfer and atomic plysics. which are bevoud the scope of this paper. | This usually requires sophisticated calculations including hydrodynamics, radiative transfer and atomic physics, which are beyond the scope of this paper. |
Instead. in the followings we use a simple analytic του-) model to estimate the fundamental parameters of SN 20024). such as Niauass and kinetic energy of the ejecta. | Instead, in the followings we use a simple analytic (“toy”-) model to estimate the fundamental parameters of SN 2002ap, such as Ni-mass and kinetic energy of the ejecta. |
The bolometric heht variation οἳ SNe can be qualitatively described by simple analytic models (Arnett.1050: Arnett. 05511 Iwamotoetal. 2002:: Maedaetal. 2003)). assuniue homologous expansion of the ejected cuvelope and (probably over-)smiplified deposition of eamuna-ravs. | The bolometric light variation of SNe can be qualitatively described by simple analytic models \cite{arnett1}; \cite{arnett2}; \cite{iwam}; \cite{maeda}) ), assuming homologous expansion of the ejected envelope and (probably over-)simplified deposition of gamma-rays. |
It is known that SNe are powered bv radioactive decay of "Ni > OCG MOF, | It is known that SNe are powered by radioactive decay of $^{56}$ Ni $\rightarrow$ $^{56}$ Co $\rightarrow$ $^{56}$ Fe. |
The basic parameter characterising the emitted cucreyv is Αν). the mass of °ONi svuthesized during tle explosion. | The basic parameter characterising the emitted energy is $M_{Ni}$, the mass of $^{56}$ Ni synthesized during the explosion. |
The guunueravs originating from radioactive decay are deposited aud thermalized in the ejecta via Couptou-scattering. | The gamma-rays originating from radioactive decay are deposited and thermalized in the ejecta via Compton-scattering. |
Because the ejecta is optically thick at carly phases. the timescale of the radiative trauster of photos to the surface is comparable to the timescale of the expansion (characterized bv Ορ the velocity of the outinost part of the ejecta). thus. the euergv remains trapped im the optically thick atmosphere. | Because the ejecta is optically thick at early phases, the timescale of the radiative transfer of photons to the surface is comparable to the timescale of the expansion (characterized by $v_{exp}$, the velocity of the outmost part of the ejecta), thus, the energy remains trapped in the optically thick atmosphere. |
As the ejecta expands. the density aud the opacity decreases; aud. the deposited cherey can escape iore and more quickly. | As the ejecta expands, the density and the opacity decreases, and the deposited energy can escape more and more quickly. |
Thus. the Πο curve has a peak at 15-20 davs past explosion. when the emitted huuinosity becomes roughly equal to the instantancous rate of the energy. deposition (Arnett.1982)). | Thus, the light curve has a peak at 15-20 days past explosion, when the emitted luminosity becomes roughly equal to the instantaneous rate of the energy deposition \cite{arnett2}) ). |
Because the trauspareucy. of the ejecta quickly decreases with tiue. the Πο variation after nani munucs the radioctive decav law and the physics of eamuna-rayv deposition. | Because the transparency of the ejecta quickly decreases with time, the light variation after maximum mimics the radioctive decay law and the physics of gamma-ray deposition. |
The very simple light curve model used in this paper is based on the assumption of free. homologous expansion of the ejecta, | The very simple light curve model used in this paper is based on the assumption of free, homologous expansion of the ejecta. |
The ejecta has a radius A(£) that increases with time as where Ry is the (ueslieible) iuitial radius of the progenitor aud (0.44 is the (coustaut) expausiou velocity. | The ejecta has a radius $R(t)$ that increases with time as where $R_0$ is the (negligible) initial radius of the progenitor and $v_{exp}$ is the (constant) expansion velocity. |
The velocity of a thin shell at fractional radius rt)R(t) las a velocity The density structure of the ejecta is assumed to be cousisting of au inner core with coustaut density. and au outer shell where the density is decreasing as power-law: where ου is the fractional radius of the core. py is the deusity in the core (1.6. fore « c9) aud à is the power-law exponent. | The velocity of a thin shell at fractional radius $x = r(t) / R(t)$ has a velocity The density structure of the ejecta is assumed to be consisting of an inner core with constant density, and an outer shell where the density is decreasing as power-law: where $x_0$ is the fractional radius of the core, $\rho_0$ is the density in the core (i.e. for $x < x_0$ ) and $n$ is the power-law exponent. |
As a consequence of the expansion. py decreases in finie as where flrg)=Gry?£3borgory?739/703.yy isa econietrie factor due to the assumed density structure. aud AM is the total mass of the ejecta, | As a consequence of the expansion, $\rho_0$ decreases in time as where $f(x_0) = ( {x_0}^3 / 3 + {x_0}^n (1- {x_0}^{3-n})/(3-n) )$ is a geometric factor due to the assumed density structure, and $M$ is the total mass of the ejecta. |
With these asstuuiptious the kinetic energy can be expressed as The energy production released as gamuna-ravs is described by the expoucutal decay law: where Ey;=3.70.10 orefe. Ec,=616-10? ere/e. Av;—13152.10"5 1, Ae,=1.0825-10.7 3. and My; is the total mass of the svuthesized Ni. | With these assumptions the kinetic energy can be expressed as The energy production released as gamma-rays is described by the exponental decay law: where $E_{Ni} = 3.70 \cdot 10^{10}$ erg/g, $E_{Co} = 6.76 \cdot 10^9$ erg/g, $\lambda_{Ni} = 1.3152 \cdot 10^{-6}$ $^{-1}$, $\lambda_{Co} = 1.0325 \cdot 10^{-7}$ $^{-1}$, and $M_{Ni}$ is the total mass of the synthesized $^{56}$ Ni. |
A fraction of the decay αιαον goes to ecucrating positrons. which are decelerated aud annihilated by the ejecta (see e.g. Capellareetal. 1997)). | A fraction of the decay energy goes to generating positrons, which are decelerated and annihilated by the ejecta (see e.g. \cite{capell}) ). |
The energy input due to positrous is approximated as Thus. the total euerev generatiou rate becomes the suu of the eamuna ray and positron cucrey productions. | The energy input due to positrons is approximated as Thus, the total energy generation rate becomes the sum of the gamma ray and positron energy productions. |
and £(/)~const.0.1 is a fairly σου approximation. | and $\xi(t) \sim \rm const. \sim 0.1$ is a fairly good approximation. |
Moreover. we assume that powers of all radiation mechanisms are dominated by Che emission from the particles 5=i. | Moreover, we assume that powers of all radiation mechanisms are dominated by the emission from the particles $\gamma = \gamma_{\rm b}$. |
This assumption is valid when (he low energv power-law index <2 aud the high energy power-law index >3 in the particle distribution. | This assumption is valid when the low energy power-law index $< 2$ and the high energy power-law index $> 3$ in the particle distribution. |
The power ratio of the inverse Compton scattering to the svnchrotron radiation is given bv where fix«l represents the IxXlein-Nishina elfect and Ü5,(/) is the energy density of the target photon field. | The power ratio of the inverse Compton scattering to the synchrotron radiation is given by where $f_{\rm KN} < 1$ represents the Klein-Nishina effect and $U_{\rm ph}(t)$ is the energy density of the target photon field. |
fy~1 for the IC/CMD and fx<0.1 for the SSC. | $f_{\rm KN} \sim 1$ for the IC/CMB and $f_{\rm KN} < 0.1$ for the SSC. |
The power of the svuchrotron radiation is given bv (e.g.Πανβίο&Lightman1979) where ο is the number of the particles around 5=>, ancl we use equation (7)). | The power of the synchrotron radiation is given by \citep[e.g.][]{rl79}
where $\gamma_{\rm b} N(\gamma_{\rm b}, t)$ is the number of the particles around $\gamma = \gamma_{\rm b}$ and we use equation \ref{eq7}) ). |
Note that although the power of the svnchrotron radiation has been conventionally compared with the spin-down power £(/). the power of the svnchrotron radiation relates {ο (he integrated spin-down energy Z0) rather than the instantaneous spin-down power L(t) in equation (24)). | Note that although the power of the synchrotron radiation has been conventionally compared with the spin-down power $L(t)$ , the power of the synchrotron radiation relates to the integrated spin-down energy $E_{\rm spin}(t)$ rather than the instantaneous spin-down power $L(t)$ in equation \ref{eq24}) ). |
Now equation (23)) becomes. in thecase of the IC/CMD μμ)= Un). | Now equation \ref{eq23}) ) becomes, in thecase of the IC/CMB $U_{\rm ph}(t) = U_{\rm CMB}$ ), |
In addition to the are itsell. other faint images are evident. especially in the subtracted image (Fig. | In addition to the arc itself, other faint images are evident, especially in the subtracted image (Fig. |
Ib). | 1b). |
Large magnilication events such as arcs are generally due (o the presence of à source near (he caustic lines (e.g.. Relsdal&Surdej(1994))). | Large magnification events such as arcs are generally due to the presence of a source near the caustic lines (e.g., \citet{Ref94}) ). |
In simple lens geometry configurations (spherical or elliptical) they are usually accompanied by additional images on the opposite side of the lens. | In simple lens geometry configurations (spherical or elliptical) they are usually accompanied by additional images on the opposite side of the lens. |
Models of the CERS03.1077 lens indicate that faint images near the position of the lensing galaxy (visible near the center of (he lens in Fig. | Models of the CFRS03.1077 lens indicate that faint images near the position of the lensing galaxy (visible near the center of the lens in Fig. |
1b) are likely (o be counter-images to the are. | 1b) are likely to be counter-images to the arc. |
ILowever. thev could also be small objects embedded within (he lensing galaxy. or simply a projection of foreground/backeround objects (with the only reservation that if they are unlensed objects. they should be then at a redshift significantly lower than that of the are source). | However, they could also be small objects embedded within the lensing galaxy, or simply a projection of foreground/background objects (with the only reservation that if they are unlensed objects, they should be then at a redshift significantly lower than that of the arc source). |
Unfortunately. we do not have color information of these images. | Unfortunately, we do not have color information of these images. |
Since their location is verv suggestive (hat they are counter images of the arc and in the absence of other data. we assume in the following that they are. in fact. lensecl images. | Since their location is very suggestive that they are counter images of the arc and in the absence of other data, we assume in the following that they are, in fact, lensed images. |
A spectrum of the lensing galaxy. CFRSO3.L077. [rom the original CFRS survey is shown in Figure 3. | A spectrum of the lensing galaxy, CFRS03.1077, from the original CFRS survey is shown in Figure 3. |
Dased mostly on the location of the bbreak ancl its overall spectral energy distribution. ΕΠΑΠΙΗΙΟΤetal.(1995) assigned a reclshilt ol z = 0.938 with a high confidence level (class 3. see LeFevreetal. (1995))). | Based mostly on the location of the break and its overall spectral energy distribution, \citet{Ham95} assigned a redshift of z = 0.938 with a high confidence level (class 3, see \citet{LeF95}) ). |
The strong bbreak and absence of any [OL] emission indicates that the galaxy is likely to be an elliptical. | The strong break and absence of any [OII] emission indicates that the galaxy is likely to be an elliptical. |
The slit was positioned EW and its width was 17775. so part of the ave was in the MOS spectrograph slit and. indeed. the spectrum of a neighboring object. ie.. (he arc. was noted bv the CFRS team but no redshift was derived. | The slit was positioned EW and its width was 75, so part of the arc was in the MOS spectrograph slit and, indeed, the spectrum of a neighboring object, i.e., the arc, was noted by the CFRS team but no redshift was derived. |
After recognizing that this was likely a lensed galaxy. this spectrum was re-reduced and analysed but still no redshift determination was possible. partly because the spectrum of the arc was overwhelmed by that of the bright ealaxy. ( | After recognizing that this was likely a lensed galaxy, this spectrum was re-reduced and analysed but still no redshift determination was possible, partly because the spectrum of the arc was overwhelmed by that of the bright galaxy. ( |
This is unlike the case of CERSI4.1311. where Cramptonetal.(1996) were able to subsequently detect strong emission lines of the background quasar). | This is unlike the case of CFRS14.1311, where \citet{Cra96} were able to subsequently detect strong emission lines of the background quasar). |
spectra of the arc alone were obtained with the OSIS spectrograph on CEIFT on 1997 August 2628. | Spectra of the arc alone were obtained with the OSIS spectrograph on CFHT on 1997 August 26–28. |
Unfortunately. the seeing was only average during these exposures (0755mm (55) and the OSIS fast guiding svstem. which would have improved the image quality somewhat. was not available. | Unfortunately, the seeing was only average during these exposures 5 -- 8) and the OSIS fast guiding system, which would have improved the image quality somewhat, was not available. |
A 07775 slit was aligned along the brightest part of the are ad PA = 20° on two of the nights. but a 1700 slit had to be used during poorer seeing on 1997 Aug 27. | A 75 slit was aligned along the brightest part of the arc at PA = $\arcdeg$ on two of the nights, but a 0 slit had to be used during poorer seeing on 1997 Aug 27. |
Two spectra were obtained on 1997 Ane 26. three on Aug 27 and five on Aug 28. | Two spectra were obtained on 1997 Aug 26, three on Aug 27 and five on Aug 28. |
All exposures were 15005. so the total exposure is equivalent to five hours. | All exposures were 1800s, so the total exposure is equivalent to five hours. |
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