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Since the pair production cross sectio rasa finction of he CM ewrey peaks at twice the electron Lass. One would exect that gamunaray photons of el0reles eus of GeV are most likely to interact wih UV ckeround plotous. | Since the pair production cross section as a function of the CM energy peaks at twice the electron mass, one would expect that gamma-ray photons of energies $\sim$ tens of GeV are most likely to interact with UV background photons. |
Thus. uusurprisiugl. modes with ie1 UV backgrounds will result i more suppression at hieh energies. | Thus, unsurprisingly, models with high UV backgrounds will result in more suppression at high energies. |
Ou the other haud. in the LDDE hoelel. while there is still more cascade radiation at ie1 energles or the Culmore model han for the Stecker hoelel. the resulting distinction between the two models is less pronunent than in the PLE nodel owing to the Meerential distribution of relatively high-flux blazars at QW redshifts. | On the other hand, in the LDDE model, while there is still more cascade radiation at high energies for the Gilmore model than for the Stecker model, the resulting distinction between the two models is less prominent than in the PLE model owing to the preferential distribution of relatively high-flux blazars at low redshifts. |
We also investigate the possibiitv that BL Lacs alc FSROs form separate populaticuns with respect to cluission and evolution bv propagating photons from distinct GLEs and SIDs for the two opulatious (Figure D). | We also investigate the possibility that BL Lacs and FSRQs form separate populations with respect to emission and evolution by propagating photons from distinct GLFs and SIDs for the two populations (Figure \ref{fig-FSRQBLL}) ). |
As both the third ECRET catalog aud the LBAS indicate that the BL Lacs tend to be more concentrated at low redshifts as compared to the FSRQs. we assume their evolution to follow the LDDE model. | As both the third EGRET catalog and the LBAS indicate that the BL Lacs tend to be more concentrated at low redshifts as compared to the FSRQs, we assume their evolution to follow the LDDE model. |
By coutrast. we asstume the FSROs to follow the PLEmodcl!®. | By contrast, we assume the FSRQs to follow the PLE. |
. For the SIDs. we asstuue the parameters deteriumed frou the LBAS for the respective sub-populations (see Section ??)). | For the SIDs, we assume the parameters determined from the LBAS for the respective sub-populations (see Section \ref{subsec-blsp}) ). |
The hard spectra of BL Lacs cause the resulting cascade radiation to be considerable relative to that of the collective intensity from iutrinsic spectra even though their collective radiation tends to be couceutrated at low redshifts iu this mocel. | The hard spectra of BL Lacs cause the resulting cascade radiation to be considerable relative to that of the collective intensity from intrinsic spectra even though their collective radiation tends to be concentrated at low redshifts in this model. |
In fact. the total intensity frou DL Lacs could dominate the EGRB at high. energies and even overproduce it. though it docs not explain the lower cherev intensity. | In fact, the total intensity from BL Lacs could dominate the EGRB at high energies and even overproduce it, though it does not explain the lower energy intensity. |
By contrast. the relatively steep spectra of FSROs are not conducive to much cascade radiation even though they are preferentially situated at | By contrast, the relatively steep spectra of FSRQs are not conducive to much cascade radiation even though they are preferentially situated at |
and heated at 1 K.min! rate (Exp. | and heated at 1 $^{-1}$ rate (Exp. |
11). | 11). |
The upper left panel in Fig. | The upper left panel in Fig. |
2 shows the CO» stretching band recorded at different temperatures during warm-up: after ice deposition at 22K, at 62 K where segregation is known to be efficient (?),, during the first ice desorption peak around 79 K, in the temperature interval between pure CO» desorption and H2O desorption, and during desorption of the trapped CO». | \ref{qms_IR} shows the $_2$ stretching band recorded at different temperatures during warm-up: after ice deposition at $22 \, \rm K$, at 62 K where segregation is known to be efficient \citep{Oberg_09_c}, during the first ice desorption peak around 79 K, in the temperature interval between pure $_2$ desorption and $_2$ O desorption, and during desorption of the trapped $_2$. |
The right panel shows the desorption rate of CO» derived from the same experiment by mass spectrometry. | The right panel shows the desorption rate of $_2$ derived from the same experiment by mass spectrometry. |
The bottom panel presents the cumulative ice loss versus temperature for this experiment, obtained both by integrating the CO» mass signal with respect to the temperature, and by integrating the CO» infrared signal recorded at specific temperatures. | The bottom panel presents the cumulative ice loss versus temperature for this experiment, obtained both by integrating the $_2$ mass signal with respect to the temperature, and by integrating the $_2$ infrared signal recorded at specific temperatures. |
The error bars on the infrared data are due to variable ice band strengths with temperature and composition. | The error bars on the infrared data are due to variable ice band strengths with temperature and composition. |
Within these uncertainties the fractional ice loss curves derived by infrared and by mass spectrometry agree well; there seems to be only a small systematic offset for the 80 — 130 K range. | Within these uncertainties the fractional ice loss curves derived by infrared and by mass spectrometry agree well; there seems to be only a small systematic offset for the 80 – 130 K range. |
This implies that the first RAIR spectrum of the ice after deposition can be used to derive quantitative results from the TPD experiments. | This implies that the first RAIR spectrum of the ice after deposition can be used to derive quantitative results from the TPD experiments. |
Figure 2 also shows that there is evidence for some ice loss between the two main desorption peaks. | Figure \ref{qms_IR} also shows that there is evidence for some ice loss between the two main desorption peaks. |
The cumulative QMS and infrared spectroscopy signals match each other at these intermediate temperatures, which points to that the measurements trace actual ice desorption in between the pure ice desorption event and the desorption of trapped volatiles. | The cumulative QMS and infrared spectroscopy signals match each other at these intermediate temperatures, which points to that the measurements trace actual ice desorption in between the pure ice desorption event and the desorption of trapped volatiles. |
The implications of this ice desorption process is discussed below, but it is important to note that this is not incorporated into the model framework and this may be a limitation to step-wise desorption models, whether using our parameterization or any of the previously published ones. | The implications of this ice desorption process is discussed below, but it is important to note that this is not incorporated into the model framework and this may be a limitation to step-wise desorption models, whether using our parameterization or any of the previously published ones. |
Quantifying this process would require an additional set of experiments where the mass spectrometer is mounted closer to the substrate to allow for the detection of very low desorption rates. | Quantifying this process would require an additional set of experiments where the mass spectrometer is mounted closer to the substrate to allow for the detection of very low desorption rates. |
Figure 3 shows the desorption of CO» from H20:CO» ice mixtures of different thicknesses (a), with different CO» concentrations (b), and heated at different rates (d). | Figure \ref{tpd_exp} shows the desorption of $_2$ from $_2$ $_2$ ice mixtures of different thicknesses (a), with different $_2$ concentrations (b), and heated at different rates (d). |
In addition there are two CO TPD curves from H35O:CO mixtures with different CO concentrations (c). | In addition there are two CO TPD curves from $_2$ O:CO mixtures with different CO concentrations (c). |
For reference, Fig. | For reference, Fig. |
3ee) presents the TPD curves of pure CO, CO», and H50 ice heated at 1 K min"!. | \ref{tpd_exp}e e) presents the TPD curves of pure CO, $_2$, and $_2$ O ice heated at 1 K $^{-1}$. |
The fraction of trapped volatile is obtained by integrating the QMS signal for temperatures above 110K and dividing it by the OMS signal integrated over the entire 20-160 K range. | The fraction of trapped volatile is obtained by integrating the QMS signal for temperatures above $110\, \rm K$ and dividing it by the QMS signal integrated over the entire 20–160 K range. |
The chosen temperature of 110 K is well below the onset of the second desorption peak and the volatiles that desorbed during the first CO? or CO desorption peak are (almost) entirely pumped, though as discussed above there seems to be a low-level type of desorption occurring between the main desorption peaks. | The chosen temperature of 110 K is well below the onset of the second desorption peak and the volatiles that desorbed during the first $_2$ or CO desorption peak are (almost) entirely pumped, though as discussed above there seems to be a low-level type of desorption occurring between the main desorption peaks. |
Whether due to finite pumping or actual desorption this results in a uncertainty in the determination of the trapped fraction,i.e.,, the choice of temperature integration limits affects the estimated amount of trapped ice by <20%. | Whether due to finite pumping or actual desorption this results in a uncertainty in the determination of the trapped fraction, the choice of temperature integration limits affects the estimated amount of trapped ice by $<$. |
. The trapped percentage of volatiles in each experiment, defined with respect to the initial volatile ice content, is reported in the second last column of Table 1. | The trapped percentage of volatiles in each experiment, defined with respect to the initial volatile ice content, is reported in the second last column of Table \ref{desorbexps}. |
The last column of Table 1 presents the trapped abundance of volatiles species with respect to the initial H5O abundance. | .The last column of Table \ref{desorbexps} presents the trapped abundance of volatiles species with respect to the initial $_2$ O abundance. |
This value is less variable compared to the trapped amount of CO/CO, with respect to the initial CO/CO)» abundance presented in the preceding column. | This value is less variable compared to the trapped amount of $_2$ with respect to the initial $_2$ abundance presented in the preceding column. |
Both Table 1 and Fig. | Both Table \ref{desorbexps} and Fig. |
3 show that for CO? and CO the percentage of trapped volatile species in the H2O ice is highly dependent on the experimental conditions; the CO, trapping fraction varies between 44 and with | \ref{tpd_exp} show that for $_2$ and CO the percentage of trapped volatile species in the $_2$ O ice is highly dependent on the experimental conditions; the $_2$ trapping fraction varies between 44 and with |
The mass in the bow shock based on the observed 60 um flux is given by (obs) = 0.042Mo(F60/135 Jy)\(D/200pe)’,, (?). | The mass in the bow shock based on the observed $60\,\mu$ m flux is given by ) = 0.042 ) )^2 , \citep{Nor97}. |
. Again, assuming the ? distance of 197 pc, and the flux of 110 Jy in 7.5 arcmin yields a shell mass of 0.033 Mo. | Again, assuming the \cite{Harp08} distance of 197 pc, and the flux of 110 Jy in 7.5 arcmin yields a shell mass of 0.033 $\msun$. |
(? derived a shell mass of 0.14 Mo, but assumed D = 400 pc for the distance to Betelgeuse.) | \citealt{Nor97} derived a shell mass of 0.14 $\msun$, but assumed $D$ = 400 pc for the distance to Betelgeuse.) |
The corresponding age of the bow shock is approximately 0000 years (see Fig. 16,, | The corresponding age of the bow shock is approximately 000 years (see Fig. \ref{fig: shellmass}, , |
solid lines). | solid lines). |
However, as discussed in Sect. 4.2,, | However, as discussed in Sect. \ref{sec: lum}, |
thisJRAS flux is most likely an overestimate, thus the age derived based on that value is an upper limit. | this flux is most likely an overestimate, thus the age derived based on that value is an upper limit. |
The shell mass from theAKARI fluxes is an order of magnitude lower (~0.0033 Μο), which would imply an even younger bow shock age. | The shell mass from the fluxes is an order of magnitude lower $\sim$$0.0033\,\msun$ ), which would imply an even younger bow shock age. |
If this is the case, however, the wind would not have sufficient time to expand to the stand-off distance. | If this is the case, however, the wind would not have sufficient time to expand to the stand-off distance. |
One possible solution is that the observed shell mass is underestimated due to uncertainties in the conversion from flux to mass (e.g. due to dust properties). | One possible solution is that the observed shell mass is underestimated due to uncertainties in the conversion from flux to mass (e.g. due to dust properties). |
In our models the shell takes ~200000 years to reach the correct Rso (see Fig. 9)), | In our models the shell takes $\sim$ 000 years to reach the correct $R_{\rm SO}$ (see Fig. \ref{fig: shockpos}) ), |
by which time the mass in the bow shock is approximately 0.02Mo. | by which time the mass in the bow shock is approximately $0.02\,\msun$. |
This is higher than the value obtained fromAKARI observations, but may be consistent within the uncertainties and with the consensus estimates for Betelgeuse's wind parameters. | This is higher than the value obtained from observations, but may be consistent within the uncertainties and with the consensus estimates for Betelgeuse's wind parameters. |
However, at this age none of our models are close to reaching a steady state. | However, at this age none of our models are close to reaching a steady state. |
Anotherpossibility is that the shell is older; however, this requires extreme stellar wind properties. | Anotherpossibility is that the shell is older; however, this requires extreme stellar wind properties. |
Utilising Eq. 3,, | Utilising Eq. \ref{eq: Mshell}, |
we see that decreasing the mass-loss rate and increasing the wind velocity will both decrease the shell mass. | we see that decreasing the mass-loss rate and increasing the wind velocity will both decrease the shell mass. |
However, the largest line widths observed in the wind are 40kmss ! Q,Fig.1), thus the wind velocity, which is usually taken to be half of this linewidth, must be approximately 20 ss! or less. | However, the largest line widths observed in the wind are $40\,$ $^{-1}$ \citep[Fig.~1]{Hug94}, thus the wind velocity, which is usually taken to be half of this linewidth, must be approximately 20 $^{-1}$ or less. |
Thus, to decrease the shell mass from a steady state value of 0.05Mo to 0.0033Μο requires that we decrease the mass-loss rate by a factor of 15, ie. M=2x1077Mo !, somewhat below the lowest estimate of ?.. | Thus, to decrease the shell mass from a steady state value of $0.05\,\msun$ to $0.0033\,\msun$ requires that we decrease the mass-loss rate by a factor of 15, i.e. $\dot{M}=2\times10^{-7} \,\msun\,$ $^{-1}$, somewhat below the lowest estimate of \citet{You93b}. |
The latter was derived from CO observations, which tends to underestimate the true mass-loss rate due to incomplete CO synthesis in Betelgeuse's stellar wind (?).. | The latter was derived from CO observations, which tends to underestimate the true mass-loss rate due to incomplete CO synthesis in Betelgeuse's stellar wind \citep{Nor97}. |
Although this lower mass-loss rate and a higher wind velocity cannot be excluded, they are outlying values in the range of observations so must be considered unlikely. | Although this lower mass-loss rate and a higher wind velocity cannot be excluded, they are outlying values in the range of observations so must be considered unlikely. |
Based on these arguments, our results suggest that Betelgeuse's bow shock is young and may not have reached a steady state yet. | Based on these arguments, our results suggest that Betelgeuse's bow shock is young and may not have reached a steady state yet. |
The circular nature and smoothness of bow shock can be naturally accounted for if this is the case, and need not be due solely to the inclination of the bow shock with respect to the plane of the sky; however, Betelgeuse's tangential velocity, V,225 ss!, and radial velocity, Vi44220.7 ss! (assuming the distance to the star is 197 pc, see Sect. 1)) | The circular nature and smoothness of bow shock can be naturally accounted for if this is the case, and need not be due solely to the inclination of the bow shock with respect to the plane of the sky; however, Betelgeuse's tangential velocity, $V_{\rm t}$ =25 $^{-1}$, and radial velocity, $V_{\rm rad}$ =20.7 $^{-1}$ (assuming the distance to the star is 197 pc, see Sect. \ref{sec: intro}) ) |
yield an inclination angle of arctan(Vi/V;44)-50?. | yield an inclination angle of $\arctan(V_{\rm t}/V_{\rm rad})$$\sim$ $^\circ$. |
Given the uncertainty in the distance to the star, this inclination angle is consistent with the ? value (56?) derived by fitting the shape of the observed arc with the analytic ? solution, even though as discussed in Sect. 4.1,, | Given the uncertainty in the distance to the star, this inclination angle is consistent with the \cite{Ueta08} value $^\circ$ ) derived by fitting the shape of the observed arc with the analytic \cite{Wil96} solution, even though as discussed in Sect. \ref{sec: morph}, |
the latter is only valid if the bow shock has reached a steady state. | the latter is only valid if the bow shock has reached a steady state. |
Utilising Betelgeuse's radial and tangential motions, the stellar velocity is (V2+V2)? ~32.5 kmss~!, the same as model B. Although the smooth appearance of the shell would seem to rule out the slow models, if the bow shock is young, the strong instabilities that characterise those simulations may not have had enough time to grow. | Utilising Betelgeuse's radial and tangential motions, the stellar velocity is $V_{\rm t}^2 +V_{\rm rad}^2)^{1/2}$$\sim$ 32.5 $^{-1}$, the same as model B. Although the smooth appearance of the shell would seem to rule out the slow models, if the bow shock is young, the strong instabilities that characterise those simulations may not have had enough time to grow. |
We can estimate the age of the bow shock by comparing the observed shape with what is predicted by the simulations in Fig. | We can estimate the age of the bow shock by comparing the observed shape with what is predicted by the simulations in Fig. |
11 [top]. | \ref{fig: shape} [top]. |
From theAKART observations, the ratio of R(0?)/R(90?) is approximately 0.7, which corresponds to an age of x30 0000 years. | From the observations, the ratio of $R(0^\circ)/R(90^\circ)$ is approximately 0.7, which corresponds to an age of $\lesssim$ 000 years. |
According to our simulations, if the bow shock is young, the bow shock tail should show strong curvature and should not be too distant from the head of the bow shock. | According to our simulations, if the bow shock is young, the bow shock tail should show strong curvature and should not be too distant from the head of the bow shock. |
Although the emission from such a tail is likely to be weak, deep, high-resolution observations may be able to detect it. | Although the emission from such a tail is likely to be weak, deep, high-resolution observations may be able to detect it. |
Indeed, there appears to be faint emission extending from the bow shock head towards the tail in theAKARI observations, e.g. the structure located at (RA,DEC) offsets of (-2,9) arcmin in the WIDE-S image of Fig. | Indeed, there appears to be faint emission extending from the bow shock head towards the tail in the observations, e.g. the structure located at (RA,DEC) offsets of (-2,9) arcmin in the WIDE-S image of Fig. |
1 of ?.. | 1 of\cite{Ueta08}. . |
A curved tail hasalready been observed for the massive O9.5 supergiant a Camelopardalis | A curved tail hasalready been observed for the massive O9.5 supergiant $\alpha$ Camelopardalis |
in the spectrum. we set the required signal-to-noise ratio to be 5c per logarithmic energy bin width of AE/E=0.1. | in the spectrum, we set the required signal-to-noise ratio to be $\sigma$ per logarithmic energy bin width of $\Delta E/E=0.1$. |
This corresponds to a 19e detection for integrated flux. | This corresponds to a $19 \sigma$ detection for integrated flux. |
From the expected flux in our standard model for M82. the required observing time to achieve this signal-to-noise ratio is 14. 19. and 330 hours for 300 GeV. | TeV. and 3 TeV. respectively. | From the expected flux in our standard model for M82, the required observing time to achieve this signal-to-noise ratio is 14, 19, and 330 hours for 300 GeV, 1 TeV, and 3 TeV, respectively. |
In the case of NGC 253. the required time is 14. 14. and 180 hours for the above 3 energy bands in our standard model. respectively. | In the case of NGC 253, the required time is 14, 14, and 180 hours for the above 3 energy bands in our standard model, respectively. |
Thus. it would be possible for CTA to obtain high energy resolution spectra to see the cascade effects within reasonable observational times. | Thus, it would be possible for CTA to obtain high energy resolution spectra to see the cascade effects within reasonable observational times. |
In this paper. we estimated the contribution from. the cascade emissions which are resulted from the inverse Compton scattering of the interstellar radiation field in the galaxies by the ee pairs created by >-ray attenuation. | In this paper, we estimated the contribution from the cascade emissions which are resulted from the inverse Compton scattering of the interstellar radiation field in the galaxies by the $e^+e^-$ pairs created by $\gamma$ -ray attenuation. |
To evaluate the >-ray attenuation. we first modeled the 5— optical depth of very high energy -ray (730 GeV) absorption by the interstellar radiation field via positron and electron. e*e. pair creation in the two nearby bright starburst galaxies. Ms2 and NGC 253. | To evaluate the $\gamma$ -ray attenuation, we first modeled the $\gamma-\gamma$ optical depth of very high energy $\gamma$ -ray $>$ 30 GeV) absorption by the interstellar radiation field via positron and electron, $e^+e^-$, pair creation in the two nearby bright starburst galaxies, M82 and NGC 253. |
To take account of the interstellar radiation. we adopted the SED model in the central star forming region provided by Siebenmorgen&Krügel(2007). | To take account of the interstellar radiation, we adopted the SED model in the central star forming region provided by \citet{sie07}. |
. Since 5-rays are created inside of the central star formation region. absorption effects from such obscured emission is important. | Since $\gamma$ -rays are created inside of the central star formation region, absorption effects from such obscured emission is important. |
We found that the attenuation. effects becomes significant above ~5 TeV and ~9 TeV for M82 and NGC 253. respectively. | We found that the attenuation effects becomes significant above $\sim 5$ TeV and $\sim9$ TeV for M82 and NGC 253, respectively. |
By using the interstellar radiation field SED model by Siebenmorgen&Krügel(2007) and our *—5 optical depth model based on it. we found that the cascade emissions would be a probe of the existence of high energy cosmic- (>10 TeV) in starburst galaxies. | By using the interstellar radiation field SED model by \citet{sie07} and our $\gamma-\gamma$ optical depth model based on it, we found that the cascade emissions would be a probe of the existence of high energy cosmic-rays $>10$ TeV) in starburst galaxies. |
The resulting total flux including cascade emission would be ~18 and ~45% | The resulting total flux including cascade emission would be $\sim$ 18 and $\sim$ |
(NVSS 1995) at A20cm. | (NVSS 1995) at $\lambda\,$ $\,$ cm. |
This may either indicate a variable radio continuum and/or the presence of the 1612 MHz OH line in Liszt's measurement. | This may either indicate a variable radio continuum and/or the presence of the $\,$ MHz OH line in Liszt's measurement. |
Here we report on observations of 3004 made with the SWS Graauw et al. | Here we report on observations of $\,$ $-$ 3004 made with the SWS $\,$ Graauw et al. |
1996) and ISOPHOT (Lemke et al. | 1996) and ISOPHOT (Lemke et al. |
1996) instruments on board ISO. together with ground-based NIR photometry. | 1996) instruments on board ISO, together with ground-based NIR photometry. |
Coverages of a 7«5.5' field centred on 3004 were made using the ISOPHOT-C 3 \ 3 pixel detector array in two spectral passbands. | Coverages of a $7'\times5.5'$ field centred on $\,$ $-$ 3004 were made using the ISOPHOT-C 3 $\times$ 3 pixel detector array in two spectral passbands. |
To avoid complete saturation of the detector the C50 (bandpass 40 - jim) and C105 (bandpass 89 - m) filters were used. | To avoid complete saturation of the detector the C50 (bandpass 40 - $\mu$ m) and C105 (bandpass 89 - $\mu$ m) filters were used. |
Since 32004. although bright. is barely discernable at A LO;m against the structured background emission from the galactic plane at the angular resolution of 45". the P32 mode was used to give the best available sky sampling. | Since $\,$ $-$ 3004, although bright, is barely discernable at $\lambda > 40\,\mu$ m against the structured background emission from the galactic plane at the angular resolution of $''$, the P32 mode was used to give the best available sky sampling. |
The data were reducec with version 7.1 of the ISOPHOT Interactive analysis package (Gabriel et al. | The data were reduced with version 7.1 of the ISOPHOT Interactive analysis package (Gabriel et al. |
1997). and further corrected for the transient behaviour of the detector. | 1997), and further corrected for the transient behaviour of the detector. |
An unresolved source was detected at the position of IRAS1793-3005 in the C50 filter with flux density 37+12 Hy at a reference wavelength of μπι (after colour correction assuming a7,[1 spectrum). | An unresolved source was detected at the position of IRAS1793-3005 in the C50 filter with flux density $37 \pm 12$ Jy at a reference wavelength of $\mu$ m (after colour correction assuming a $\nu^{+4}$ spectrum). |
À (confusior limited) upper limit of JJy was obtained in the C105 filter. | A (confusion limited) upper limit of Jy was obtained in the C105 filter. |
In addition. the source was observed with the ESO-MPG 2.2m telescope on La Silla in June 1998 using the IRAC2B camera in the J. H. and K bands. | In addition, the source was observed with the ESO-MPG 2.2m telescope on La Silla in June 1998 using the IRAC2B camera in the J, H, and K bands. |
Applying the data reduction techniques described in Philipp et al. ( | Applying the data reduction techniques described in Philipp et al. ( |
1999). we derived the flux densities given in Table 2.. | 1999), we derived the flux densities given in Table \ref{tabnirflux}. |
As a follow-up to the detection of 3004 as such à prominent source in the scans in 7 with ISOPHOT. a full spectral sean was made on 20/ February 1998 in the j5jm spectral range using the SWS instrument in AOTOI mode with scanning speed 3. | As a follow-up to the detection of $\,$ $-$ 3004 as such a prominent source in the scans in $l$ with ISOPHOT, a full spectral scan was made on $^{th}$ February 1998 in the $\,$ $\,$ $\mu$ m spectral range using the SWS instrument in AOT01 mode with scanning speed 3. |
The pointing was specified directly towards the probable radio counterpart measured by Liszt (1992) at (BI950) 17397"22.7* —30*01'16.0". thus ensuring that the full 15" 10" extent of the radio source (at 13°) would be enclosed by the SWS apertures for the position angle of the observation. | The pointing was specified directly towards the probable radio counterpart measured by Liszt (1992) at (B1950) $17^h39^m22.7^s$ $-30^\circ04'16.0''$, thus ensuring that the full $''\,$ $\,$ $''$ extent of the radio source (at $\,$ $^\circ\,$ ) would be enclosed by the SWS apertures for the position angle of the observation. |
The spectrum was fully reduced with LA and OSIA (version of November 257 1998 and Version 1.0 respectively: Thronley et al. | The spectrum was fully reduced with IA and OSIA (version of November $^{th}$ 1998 and Version 1.0 respectively; Thronley et al. |
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