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In this plane. a higher Iddington ratio corresponds to a line with smaller positive slope.
In this plane, a higher Eddington ratio corresponds to a line with smaller positive slope.
We have avoided binning objects by their Eddington ratio in thispaper. simply because the shapes of those bins would not lend themselves to easy comparisons.
We have avoided binning objects by their Eddington ratio in thispaper, simply because the shapes of those bins would not lend themselves to easy comparisons.
We would.
We would,
carbon abundance predicted by stellar models gave rise to only a tiny change in the EW.
carbon abundance predicted by stellar models gave rise to only a tiny change in the EW.
Therefore we applied the maximum decrease in the carbon to oxygen ratio found by Erbetal.(2010) at z~2 so as to investigate the variation of the initial carbon/oxygen abundance ratio.
Therefore we applied the maximum decrease in the carbon to oxygen ratio found by \citet{erb} at $z\sim2$ so as to investigate the variation of the initial carbon/oxygen abundance ratio.
We present two sets of models, one with normal Solar-scaled carbon abundances and a second with the carbon abundance decreased by a factor of four.
We present two sets of models, one with normal Solar-scaled carbon abundances and a second with the carbon abundance decreased by a factor of four.
We only perform this decrease for main-sequence OB stars when we use the OB model atmospheres of Smith,Norris&Crowther (2002).
We only perform this decrease for main-sequence OB stars when we use the OB model atmospheres of \citet{crow}.
. The result is general reduction in the strength of the absorption line as measured by the equivalent width.
The result is general reduction in the strength of the absorption line as measured by the equivalent width.
As Figure 1 demonstrates this reduction can go to explain some of the scatter in the observed EWs of in nearby galaxies, although the majority of the scatter is due to variations in their star formation history.
As Figure \ref{civew} demonstrates this reduction can go to explain some of the scatter in the observed EWs of in nearby galaxies, although the majority of the scatter is due to variations in their star formation history.
The procedure outlined above yields a synthetic spectrum appropriate to each time-step of a stellar evolution model.
The procedure outlined above yields a synthetic spectrum appropriate to each time-step of a stellar evolution model.
We can then combine the spectra for each star together to produce the integrated spectrum for a synthetic stellar population.
We can then combine the spectra for each star together to produce the integrated spectrum for a synthetic stellar population.
To do this we use the initial mass function described by Kroupa(2002).
To do this we use the initial mass function described by \citet{kroupa2002}.
. This uses an IMF power-law slope of -1.3 between 0.1 and MMo, and a slope of -2.35 from 0.5 to MMs.
This uses an IMF power-law slope of -1.3 between 0.1 and $_{\odot}$, and a slope of -2.35 from 0.5 to $_{\odot}$.
Finally in our spectral synthesis we include the contribution from nebular emission.
Finally in our spectral synthesis we include the contribution from nebular emission.
In star-forming galaxies, interstellar gas is ionised by the stellar continuum emitted blueward of912A,, and upon recombination it emits a nebular continuum.
In star-forming galaxies, interstellar gas is ionised by the stellar continuum emitted blueward of, and upon recombination it emits a nebular continuum.
Neglecting this emission would lead to an incorrect estimate of the equivalent widths of emission lines and incorrect broad-band colours (Zackrisson,Bergvall&Leitet2008;Mollaetal.
Neglecting this emission would lead to an incorrect estimate of the equivalent widths of emission lines and incorrect broad-band colours \citep{zack,molla}.
2009).. We use the radiative transfer program (Ferlandetal.1998) to produce a detailed model of the output nebular emission spectrum excited by our stellar spectra.
We use the radiative transfer program \citep{cloudy} to produce a detailed model of the output nebular emission spectrum excited by our stellar spectra.
The model output is sensitive to the chosen geometry, inner radius and composition of the gas used in the code.
The model output is sensitive to the chosen geometry, inner radius and composition of the gas used in the code.
The details of our illustrative nebular emission model are identical to those in Eldridge&Stan-way (2009), and we output the final continuum and line strengths for use in our synthetic spectra.
The details of our illustrative nebular emission model are identical to those in \citet{ES09}, and we output the final continuum and line strengths for use in our synthetic spectra.
We consider in this work the and spectral lines, both of which exhibit a broad component whose strength is determined by the stellar spectrum.
We consider in this work the and spectral lines, both of which exhibit a broad component whose strength is determined by the stellar spectrum.
We have already shown
We have already shown
A transit of GJ1I2I4b was observed in Sloan r-band (1.2627 nm) on May 26. 2010 between 2:57 UT and 5:15 UT with the Wide Field Camera (WFC) on the 2.5 meter Isaac Newton Telescope (INT).
A transit of GJ1214b was observed in Sloan r-band $\lambda_c$ =627 nm) on May 26, 2010 between 2:57 UT and 5:15 UT with the Wide Field Camera (WFC) on the 2.5 meter Isaac Newton Telescope (INT).
An exposure time of 60 seconds was used resulting in 89 frames with an average cycle time of 93 seconds.
An exposure time of 60 seconds was used resulting in 89 frames with an average cycle time of 93 seconds.
Only the central detector of the WFC (CCD4) was used for the analysis. with a pixel scale of 0.33°7/pixel. this CCD has a field of view of 675" by 13507. sufficient to observe a number of reference stars of similar brightness simultaneously with the target.
Only the central detector of the WFC (CCD4) was used for the analysis, with a pixel scale of 0.33”/pixel, this CCD has a field of view of 675” by 1350”, sufficient to observe a number of reference stars of similar brightness simultaneously with the target.
The moon was almost full and the conditions were strongly non-photometric. with the transparency dropping to below during several frames (see top panel of Fig. 1)).
The moon was almost full and the conditions were strongly non-photometric, with the transparency dropping to below during several frames (see top panel of Fig. \ref{fig:LC_raw}) ).
On July 29. 2010. a transit of GJ214. was observed in I- 61.2822 nm) with the same instrument.
On July 29, 2010, a transit of GJ1214 was observed in I-band $\lambda_c$ =822 nm) with the same instrument.
The observations started at 21:23 UT and lasted for just over 3 hours.
The observations started at 21:23 UT and lasted for just over 3 hours.
An exposure time of 50 seconds was used. resulting in a total of 142 frames with an average cycle time of 81 seconds.
An exposure time of 50 seconds was used, resulting in a total of 142 frames with an average cycle time of 81 seconds.
In this case the night was photometric.
In this case the night was photometric.
Since GJI214b ts about 5.5 times brighter in I-band than in r-band. we significantly defocused the telescope in the I-band in order to keep the peak count-levels in the linear regime of the detector.
Since GJ1214b is about 5.5 times brighter in I-band than in r-band, we significantly defocused the telescope in the I-band in order to keep the peak count-levels in the linear regime of the detector.
This has the added benefit that the light is spread over more pixels. reducing the impact of flat-fielding errors.
This has the added benefit that the light is spread over more pixels, reducing the impact of flat-fielding errors.
On the night of July 3. 2010. we obtained simultaneous observations of GJ1214b in the g 61,2459 nm). r 61.2622 nm). i Gl.2764 nm). and z-band 61,2899 nm) with the GROND instrument (?) on the 2.2 meter MPI/ESO telescope at La Silla in Chile.
On the night of July 3, 2010, we obtained simultaneous observations of GJ1214b in the g $\lambda_c$ =459 nm), r $\lambda_c$ =622 nm), i $\lambda_c$ =764 nm), and z-band $\lambda_c$ =899 nm) with the GROND instrument \citep{greineretal08} on the 2.2 meter MPI/ESO telescope at La Silla in Chile.
The field of view in each of the wavelength channels is 5.4. by 5.4. which ts sufficient to observe both GJI214 and a set of reference stars simultaneously.
The field of view in each of the wavelength channels is 5.4' by 5.4', which is sufficient to observe both GJ1214 and a set of reference stars simultaneously.
The observations started at 00:16 UT and lasted until 04:06 UT.
The observations started at 00:16 UT and lasted until 04:06 UT.
During this time we obtained 280 frames in each of the four optical bands.
During this time we obtained 280 frames in each of the four optical bands.
The exposure time was varied from 20 to 30 seconds to avoid saturation of the CCDs.
The exposure time was varied from 20 to 30 seconds to avoid saturation of the CCDs.
The average cycle time was 50 seconds.
The average cycle time was 50 seconds.
We obtained a K, band transit observations with. the NOTCam instrument on the Nordic Optical Telescope (NOT) simultaneously with our INT r-band observations on May 26. 2010.
We obtained a $_s$ band transit observations with the NOTCam instrument on the Nordic Optical Telescope (NOT) simultaneously with our INT r-band observations on May 26, 2010.
The observations were carried out in service mode. and the wide field imaging opties and a K,-band filter (1.22.15 um) were used.
The observations were carried out in service mode, and the wide field imaging optics and a $_s$ -band filter $\lambda_c$ =2.15 $\mu$ m) were used.
The pixelscale of this setup is 0.234pixel. resulting in a field of view of the detector of 4 by 4 areminutes.
The pixelscale of this setup is 0.234”/pixel, resulting in a field of view of the detector of 4 by 4 arcminutes.
This field of view is sufficient to allow simultaneous observations of one reference star of similar brightness to GJI214 as well as a reference star that is 4x fainter than GJI214.
This field of view is sufficient to allow simultaneous observations of one reference star of similar brightness to GJ1214 as well as a reference star that is $\times$ fainter than GJ1214.
The field of view of the detector was rotated to make sure that bad regions on the detector were avoided for all three stars.
The field of view of the detector was rotated to make sure that bad regions on the detector were avoided for all three stars.
Since the NOT is located on the same mountain as the INT. these observations suffer from the same non-photometric conditions. with the transparency dropping to for parts of the light curve (see Fig. 1)).
Since the NOT is located on the same mountain as the INT, these observations suffer from the same non-photometric conditions, with the transparency dropping to for parts of the light curve (see Fig. \ref{fig:LC_raw}) ).
This strongly affects the observations. since the sky background dominates over the object flux for the larger apertures. especially during times of low transparency.
This strongly affects the observations, since the sky background dominates over the object flux for the larger apertures, especially during times of low transparency.
Since GJI214 is bright at near-infrared wavelengths. we defocused the telescope in order to allow for the relatively long integration time and reducing the sensitivity to flat-fielding errors. although this also increases the impact of the sky background.
Since GJ1214 is bright at near-infrared wavelengths, we defocused the telescope in order to allow for the relatively long integration time and reducing the sensitivity to flat-fielding errors, although this also increases the impact of the sky background.
The exposure time was set to 4 seconds. to allow for relatively efficient. observations. with the large overheads induced by the NOTCam system.
The exposure time was set to 4 seconds, to allow for relatively efficient observations, with the large overheads induced by the NOTCam system.
The average cycle time was 16 seconds. allowing us to capture 518 frames in 2 hours and 25 minutes of observations.
The average cycle time was 16 seconds, allowing us to capture 518 frames in 2 hours and 25 minutes of observations.
In order to increase the stability of the system for the observations. and to decrease the telescope overheads. we observed in staring mode. with guiding. keeping the centroid of the star constant to within 4 pixels during the observations.
In order to increase the stability of the system for the observations, and to decrease the telescope overheads, we observed in staring mode, with guiding, keeping the centroid of the star constant to within 4 pixels during the observations.
Since this observation strategy does not allow us to subtract the background from the images. we obtained a set of dithered observations after our transit observation. from which we constructed a background map.
Since this observation strategy does not allow us to subtract the background from the images, we obtained a set of dithered observations after our transit observation, from which we constructed a background map.
ensing were recently studied in detail by ?.BAIL using clusters formed in a cosmological simulation independent of he MS run at. somewhat lower resolution.
lensing were recently studied in detail by \citet[BK11]{Becker_Kravtsov_2011} using clusters formed in a cosmological simulation independent of the MS run at somewhat lower resolution.
In this study. he mass derived from WL was found to be biased. low ato a level of ~5'A.
In this study, the mass derived from WL was found to be biased low at a level of $\sim -5\%$.
Considering that. these authors emploved a slightly dillerent reconstruction method in which xckeround. galaxies were used over a radial range [rom 1’ o 20. from the cluster centre and a shear profile formed. rom them. our results are in good quantitative agrecmicnt (sce also Fig.
Considering that these authors employed a slightly different reconstruction method in which background galaxies were used over a radial range from $^\prime$ to $^\prime$ from the cluster centre and a shear profile formed from them, our results are in good quantitative agreement (see also Fig.
B83 in the appendix. where we analyse our simulation using the same radial range as DIX11).
\ref{fig:cuttest} in the appendix, where we analyse our simulation using the same radial range as BK11).
The scatter determined by DINXI1 (~20%) is very close to the evel we derive (~ 25%): we note. however. that these two numbers were derived: using two slightly. dillerent. analysis methocls.
The scatter determined by BK11 $\sim 20\%$ ) is very close to the level we derive $\sim 25\%$ ); we note, however, that these two numbers were derived using two slightly different analysis methods.
Biases in both mass and concentration have also been stucied by 2.OLILLI.. who found a mass bias similar to that in? and presented here.
Biases in both mass and concentration have also been studied by \citet[OH11]{Oguri_Hamana_2011}, who found a mass bias similar to that in \citet{Becker_Kravtsov_2011} and presented here.
However. in contrast to our results. they find a very laree. positive concentration bias of vs20%.
However, in contrast to our results, they find a very large, positive concentration bias of $\sim 20\%$.
]t is presently unclear what the origin of this dillerence is.
It is presently unclear what the origin of this difference is.
We speculate that it may originate from a dillerence in the weak lensing simulation method. between our two stuclies: while our results are based. on direct. fitting of a high resolution weak lensing simulation. OLLI] analyse an analytic circularly svmmoetric shear profile that was derived by stacking5 (mock) rav-traced lensing5 observations of 5galaxy clusters formed in à low resolution cosmological simulation.
We speculate that it may originate from a difference in the weak lensing simulation method between our two studies: while our results are based on direct fitting of a high resolution weak lensing simulation, OH11 analyse an analytic circularly symmetric shear profile that was derived by stacking (mock) ray-traced lensing observations of galaxy clusters formed in a low resolution cosmological simulation.
The azimuthal averagingDIA is expected to smooth out. the presence of substructure and triaxiality. both of which tend to bias the concentration low (see 5.3.3)).
The azimuthal averaging is expected to smooth out the presence of substructure and triaxiality, both of which tend to bias the concentration low (see \ref{sec:biasorigin}) ).
Having established the extent of the scatter and bias in WL reconstructions of cluster haloes. we now aim to find physical explanations for them.
Having established the extent of the scatter and bias in WL reconstructions of cluster haloes, we now aim to find physical explanations for them.
This is an interesting question in its own right. but might also allow an identification of possible ways to reduce these systematic errors.
This is an interesting question in its own right, but might also allow an identification of possible ways to reduce these systematic errors.
Any potential error sources can be broadly grouped into two categories: Those due to the background galaxies used in the reconstruction (ic. their unknown intrinsic ellipticities and finite number. in general also their intrinsic alignment due to cosmic shear). and those due to the cluster itself. such as halo triaxiality and substructure.
Any potential error sources can be broadly grouped into two categories: Those due to the background galaxies used in the reconstruction (i.e., their unknown intrinsic ellipticities and finite number, in general also their intrinsic alignment due to cosmic shear), and those due to the cluster itself, such as halo triaxiality and substructure.
In this section we show that. in the case of large statistical samples of clusters such as we have studied here and those to be derived. [rom the DES and LAST. the latter is dominated by the former only [or clusters with masses below a few 107M..
In this section we show that, in the case of large statistical samples of clusters such as we have studied here and those to be derived from the DES and LSST, the latter is dominated by the former only for clusters with masses below a few $10^{14} {\rm M}_\odot$.
Our strategy for assessing the importance of these various error contributions involves making two acelitional reconstructions of our cluster sample. designed. to. bridge the eap between the WL analysis based on particles within a lO h! Alpe box on the one side and the 3D fitting procedure within a radius ουυ on the other.
Our strategy for assessing the importance of these various error contributions involves making two additional reconstructions of our cluster sample, designed to bridge the gap between the WL analysis based on particles within a 10 $h^{-1}$ Mpc box on the one side and the 3D fitting procedure within a radius $r_{200}$ on the other.
In. the first of these. which we will refer to as “perfect WL. we use a very high density (η=3002 arcmin 7) of perfectly. circular background: galaxies πιο. o= 0.0). which eliminates the influence of shape noise and — essentially. we are now analysing the (reduced) shear field. g clirectly.
In the first of these, which we will refer to as “perfect WL”, we use a very high density $n = 300$ $^{-2}$ ) of perfectly circular background galaxies (i.e., $\sigma = 0.0$ ), which eliminates the influence of shape noise and — essentially, we are now analysing the (reduced) shear field $g$ directly.
In the second method. we approach the 3D fit even further by constructing the convergence field 5s and thus the shear only from those particles that lie within a (3D) distance of ους Crom the cluster centre. the same set of particles upon which the 3D fit is based.
In the second method, we approach the 3D fit even further by constructing the convergence field $\kappa$ — and thus the shear — only from those particles that lie within a (3D) distance of $r_{200}$ from the cluster centre, the same set of particles upon which the 3D fit is based.
We will refer to this method as "spherical WL.
We will refer to this method as “spherical WL”.
One problem with this approach is that the expression of ?. for g assumes à matter distribution extending to infinity. so fitting il to à catalogue of galaxy distortions based only on the matter distribution inside roug alone. which necessarily contains less total mass. will result in severe biases in both mass and.
One problem with this approach is that the expression of \citet{Wright_Brainerd_2000} for $g$ assumes a matter distribution extending to infinity, so fitting it to a catalogue of galaxy distortions based only on the matter distribution inside $r_{200}$ alone, which necessarily contains less total mass, will result in severe biases in both mass and.
. We overcome this by instead fitting the reduced shear fron an NEW profile that is. like our data. truncated: at σου (?: ?)).
We overcome this by instead fitting the reduced shear from an NFW profile that is, like our data, truncated at $r_{200}$ \citealt{Takada_Jain_2003a}; \citealt{Takada_Jain_2003b}) ).
We note that it is straightforward to also fit the projected mass profile of 7. directly to the projected matter within ους.
We note that it is straightforward to also fit the projected mass profile of \citet{Takada_Jain_2003a} directly to the projected matter within $r_{200}$.
This method approaches the 3D fit even closer. with the only remaining dillerence being a fit in 2D vs. one in 3D. We have done this. ancl found that the results are in very close agreement with those of the "spherical. WL' methocl.
This method approaches the 3D fit even closer, with the only remaining difference being a fit in 2D vs. one in 3D. We have done this, and found that the results are in very close agreement with those of the `spherical WL' method.
The bias and scatter as defined. in equation in masses and concentrations resulting [rom both these methocs are shown in Fig.
The bias and scatter as defined in equation in masses and concentrations resulting from both these methods are shown in Fig.
4 and 5. respectively: the perfect WL fit represented by red lines. the spherical WL fit by blue ones.
\ref{medians} and \ref{fig:scatter} respectively, the perfect WL fit represented by red lines, the spherical WL fit by blue ones.
For ease of comparison. we also include the values determined from our default WL simulation as found in ‘Table 1 as black lines.
For ease of comparison, we also include the values determined from our default WL simulation as found in Table \ref{table} as black lines.
We can judgee the influence of errors. associated with the background. galaxies by comparing the "default and “perfect” WL reconstructions (black and red. curves. in Fies.
We can judge the influence of errors associated with the background galaxies by comparing the “default” and “perfect” WL reconstructions (black and red curves in Figs.
5 4 and 5)). the latter one. as explained1 above. not being allected by them.
\ref{medians} and \ref{fig:scatter}) ), the latter one, as explained above, not being affected by them.
Focusing first on scatter. in. default) vs. perfect WL (Fig. 5)).
Focusing first on scatter in default vs. perfect WL (Fig. \ref{fig:scatter}) ),
we find a very similar picture for both mass and concentration: it is comparable for the highest-mass clusters. but while the latter is almost. mass-independoent. the former increases considerably with decreasing. cluster mass.
we find a very similar picture for both mass and concentration: it is comparable for the highest-mass clusters, but while the latter is almost mass-independent, the former increases considerably with decreasing cluster mass.
This is what would be expected. from the influence of shape noise: Less massive haloes. producing a weaker shear signal. viel a lower signal-to-noise ratio than their high-mass counterparts: in the total absence of shape noise. however. the decreasing shear signal is irrelevant.
This is what would be expected from the influence of shape noise: Less massive haloes, producing a weaker shear signal, yield a lower signal-to-noise ratio than their high-mass counterparts; in the total absence of shape noise, however, the decreasing shear signal is irrelevant.
As with scatter. the variation in bias between default. and
As with scatter, the variation in bias between default and
Additional evidence that we are distinguishiug regions domiuated by AGN activity from H I 'eejous iu NGCLILO comes [rom the line widths aud the lack of stroug contiuuuim in tle H II sources.
Additional evidence that we are distinguishing regions dominated by AGN activity from H II regions in NGC4410 comes from the line widths and the lack of strong continuum in the H II sources.
The nucleus of NGC [110A has relatively broad lines (FWHM ~ 600 kin 1j
The nucleus of NGC 4410A has relatively broad lines (FWHM $\sim$ 600 km $^{-1}$ ).
The NGC 10Η iucleus aud the filamentary features also appear to have somewhat broad lines (FWHM ~ LOO ans 1) compared to those of the knots(FWHM ~ 200 — 300 kins 1). which are only marginally 'esolved.
The NGC 4410B nucleus and the filamentary features also appear to have somewhat broad lines (FWHM $\sim$ 400 km $^{-1}$ ) compared to those of the knots(FWHM $\sim$ 200 $-$ 300 km $^{-1}$ ), which are only marginally resolved.
The knot sources have very little continuum. cousistent with their identification as H HH 'eglons.
The knot sources have very little continuum, consistent with their identification as H II regions.
The velocity structure of the ionized gas iu this system (Figure 2)) is intriguing.
The velocity structure of the ionized gas in this system (Figure \ref{fig7}) ) is intriguing.
The two iuclei. the northwestern filament. aud the extended emission to the east of the NGC [110A nucleus are considerably redshifted (7110 ki | — 7500 kins 1) from the H LE regious and the extended emission to the southwest and east of the NGC 1110 nucleus. centered at 7130 — 7310 kin |).
The two nuclei, the northwestern filament, and the extended emission to the east of the NGC 4410A nucleus are considerably redshifted (7440 km $^{-1}$ $-$ 7500 km $^{-1}$ ) from the H II regions and the extended emission to the southwest and east of the NGC 4410A nucleus, centered at 7130 $-$ 7340 km $^{-1}$ ).
HH is possible that we are viewing line emission [rom two interacting systems with a velocity separation of ~200 km +. or that the H II regions were originally associated with one of the nuclei.
It is possible that we are viewing line emission from two interacting systems with a velocity separation of $\sim 200$ km $^{-1}$, or that the H II regions were originally associated with one of the nuclei.
The IUE data were extracted and. calibrated by the IUE New Spectral Imaging Processing System (NEWSIPS) software in two ways: the staudard extraction method which is optimized lor a poiut source. aud a re-extraction aud reprocessing using a method optimized for an extended source.
The IUE data were extracted and calibrated by the IUE New Spectral Imaging Processing System (NEWSIPS) software in two ways: the standard extraction method which is optimized for a point source, and a re-extraction and reprocessing using a method optimized for an extended source.
The two SWP spectra are plotted in Figure ) 3..
The two SWP spectra are plotted in Figure \ref{fig8}. .
The 1216A ἵνα feature is strongly detected in this source. with a total line lux of 1.0 + 0.2 x P erg tem 7.
The ${\rm \AA}$ $\alpha$ feature is strongly detected in this source, with a total line flux of 1.0 $\pm$ 0.2 $\times$ $^{-13}$ erg $^{-1}$ $^{-2}$.
A possible 15184-1551A. C IV feature may be present at the level of 6 + ‘ 9 −≻≍⋯⊔↩⋅∑≟⊳∖↓∢∙∐↕−⋅∣≻⋯↕∐⊳∖↥≺↵⋜↕⊓⊔⋅≺↵∢∙∩∐∐∙∐⇂≺↵⊳∖∖∖↽∐∐↕∐≺↵↥↽≻↓⋅≺↵⊳∖≺↵∐∢∙≺↵∩↥⊳∖∩⋯≺↵∢∙∩⊳∖∐∐∢∙↕⋅⋜↕⊽∖⊽⊳∖⋜↕∐≺⇂ ⋅ ⋅ ⋅ ⋅ ⋅ ⋅ the wavelength centroid of the feature places it significantly bluer (by about. £00 kin 1) of the centroid of the Lya feature. and lies about £5” south of the position of the continuum source.
A possible ${\rm \AA}$ C IV feature may be present at the level of 6 $\pm$ 2 $\times$ $^{-14}$ erg $^{-1}$ $^{-2}$, but this feature coincides with the presence of some cosmic rays and the wavelength centroid of the feature places it significantly bluer (by about 400 km $^{-1}$ ) of the centroid of the $\alpha$ feature, and lies about $''$ south of the position of the continuum source.
If real. it may be associated with Ixuot #2. but we are doubtful.
If real, it may be associated with Knot $\#$ 2, but we are doubtful.
The mean contiuuum between 1350-1150 observed frame in the SWP spectrum is Peres 7 A1,
The mean continuum between 1350-1450 observed frame in the SWP spectrum is $1.1\pm0.2 \times 10^{-15}$ erg $^{-1}$ $^{-2}$ $^{-1}$.
The contimuun flux from the SWP spectrum obtained with the extended source algoritlun is statistically equal to that fouud with the point source algoritlin centered ou the coutinuumn peak (1.2c0.2x1015 eres tem7 At). indicating that the continuum enmission ix unresolved by IUE (FWHM < 3755).
The continuum flux from the SWP spectrum obtained with the extended source algorithm is statistically equal to that found with the point source algorithm centered on the continuum peak $1.2\pm0.2 \times 10^{-15}$ erg $^{-1}$ $^{-2}$ $^{-1}$ ), indicating that the continuum emission is unresolved by IUE (FWHM $<$ 5).
The LWP spectrum (uot shown) contains only faint. [lat continuum emission. at a mean level of b+ 2 x Meres Fem? Lat30004.
The LWP spectrum (not shown) contains only faint, flat continuum emission, at a mean level of 4 $\pm$ 2 $\times$ $^{-16}$ erg $^{-1}$ $^{-2}$ $^{-1}$ at.
Lya is clearly visible in both extracted spectra. and the fIux obtained with the extended source methocl is higher than that derived asstunineg a polut source.
$\alpha$ is clearly visible in both extracted spectra, and the flux obtained with the extended source method is higher than that derived assuming a point source.
In Figure 3.. the amplituce of the Lya emission line [rom the extendedsource extraction (71.210. theres Fem 7 4) is higher than that obtained from the point-source extraction (~0.5x10 theres tem 7 +). indicating that Lya is extendec.
In Figure \ref{fig8}, , the amplitude of the $\alpha$ emission line from the extendedsource extraction $\sim 1.2 \times 10^{-14}$ erg $^{-1}$ $^{-2}$ $^{-1}$ ) is higher than that obtained from the point-source extraction $\sim 0.5 \times 10^{-14}$ erg $^{-1}$ $^{-2}$ $^{-1}$ ), indicating that $\alpha$ is extended.
The extent of the Lya emission is >20". since the emission seems to fill the IUE aperture.
The extent of the $\alpha$ emission is $>20''$, since the emission seems to fill the IUE aperture.
candles in cosmology.
candles in cosmology.
Finally. we are very interested iu exploring the observational consequences of tlhe AITT in more detail.
Finally, we are very interested in exploring the observational consequences of the MTI in more detail.