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?— found correlations and anti-correlations between near-infrared extinction and molecular clouds traced by CO emission could be used to resolve the ambiguity towards molecular clouds.
\citet{dame1991} found correlations and anti-correlations between near-infrared extinction and molecular clouds traced by CO emission could be used to resolve the ambiguity towards molecular clouds.
This method worked by assuming that the anti-correlation was the result of dust associated with foreground clouds.
This method worked by assuming that the anti-correlation was the result of dust associated with foreground clouds.
?— used a statistica approach to derive a luminosity-physical diameter correlation for HII regions.
\citet{paladini2004} used a statistical approach to derive a luminosity-physical diameter correlation for HII regions. \citet{downes1980},
2.. 22 and ? have used a combination of 10a radio recombination lines (RRLs) and formaldehyde (H»CO) absorption measurements to assign near and far solutions for HII regions.
\citet{araya2001, araya2002} and \citet{sewilo2004} have used a combination of $\alpha$ radio recombination lines (RRLs) and formaldehyde $_2$ CO) absorption measurements to assign near and far solutions for HII regions.
Although these studies have had reasonable success in resolving ambiguities they tend to be used only for a particular type of object and/or require a specific set of conditions and so are not universally applicable.
Although these studies have had reasonable success in resolving ambiguities they tend to be used only for a particular type of object and/or require a specific set of conditions and so are not universally applicable.
Another technique that has been applied successfully to both UCHII regions and YSOs combines cem HI absorption with a molecular tracer such as CO and implicitly assumes that the source is still embedded within its natal molecular cloud.
Another technique that has been applied successfully to both UCHII regions and YSOs combines cm HI absorption with a molecular tracer such as $^{13}$ CO and implicitly assumes that the source is still embedded within its natal molecular cloud.
For HII regions this technique relies on the principle that if a continuum source is located at the near distance. the HI data will not show any absorption from clouds with velocities between that of the HII region and the tangent point. since these are located behind the continuum source with respect to our line of sight.
For HII regions this technique relies on the principle that if a continuum source is located at the near distance, the HI data will not show any absorption from clouds with velocities between that of the HII region and the tangent point, since these are located behind the continuum source with respect to our line of sight.
If absorption is present in the HI data at distances between the continuum sources and the velocity of the tangent point the source must be located at the far distance (e.g. 2?)
If absorption is present in the HI data at distances between the continuum sources and the velocity of the tangent point the source must be located at the far distance (e.g., \citealt{kuchar1994,kolpak2003}) ).
For YSOs this technique looks for HI self-absorption (SA) at the same velocity as the source and works on the principle that if the sources host cloud is located at the near distance. it would lie in front of a significant column of warmer HI. resulting in absorption by cold HI associated with the cloud.
For YSOs this technique looks for HI self-absorption (SA) at the same velocity as the source and works on the principle that if the source's host cloud is located at the near distance, it would lie in front of a significant column of warmer HI, resulting in absorption by cold HI associated with the cloud.
Conversely. the absence of an absorption dip would imply the source is located at the far distance (e.g.. 22).
Conversely, the absence of an absorption dip would imply the source is located at the far distance (e.g., \citealt{jackson2002,busfield2006}) ).
In this study we will focus on the sub-sample of HHII regions and MYSO candidates located within the GRS longitude range.
In this study we will focus on the sub-sample of HII regions and MYSO candidates located within the GRS longitude range.
Combining these two data sets has the potential to solve the distance ambiguities towards approximately one-third of inner-Galaxy sources. and 75 per cent of the sources located in the Northern Galactic Plane.
Combining these two data sets has the potential to solve the distance ambiguities towards approximately one-third of inner-Galaxy sources, and 75 per cent of the sources located in the Northern Galactic Plane.
The distance ambiguities in this region have been solved using a combination of the two HI absorption techniques described in the previous paragraph for a flux-limited sample of molecular clouds (2)..
The distance ambiguities in this region have been solved using a combination of the two HI absorption techniques described in the previous paragraph for a flux-limited sample of molecular clouds \citep{roman2009}.
Both methods have been used for molecular clouds that are also found to be associated with a HIT region with the results of both techniques being check to ensure consistency between the two.
Both methods have been used for molecular clouds that are also found to be associated with a HII region with the results of both techniques being check to ensure consistency between the two.
The distance solutions have been checked with many of the previous surveys and the results are found to be in reasonable agreement.
The distance solutions have been checked with many of the previous surveys and the results are found to be in reasonable agreement.
Therefore the GRS catalogue provides robust distances to a complete sample of molecular clouds located within the first quadrant of the Galactic plane.
Therefore the GRS catalogue provides robust distances to a complete sample of molecular clouds located within the first quadrant of the Galactic plane.
In the Northern Galactic Plane we have 427 RMS sources located within the solar circle that are atfected by these kinematic distance ambiguities.
In the Northern Galactic Plane we have 427 RMS sources located within the solar circle that are affected by these kinematic distance ambiguities.
Searching this sample we have identified 306 young massive stars HHIIs and MYSO candidates) located. within the GRS region.
Searching this sample we have identified 306 young massive stars HIIs and MYSO candidates) located within the GRS region.
In order to obtain kinematic velocities for these. we extracted. spectra from the GRS data cubes and fitted them with Gaussian profiles.
In order to obtain kinematic velocities for these, we extracted spectra from the GRS data cubes and fitted them with Gaussian profiles.
In cases where more than one significant emission peak was present in the CO spectrum. we have used archival maser or high-density gas tracers to determine the source velocity (see ?. for more details).
In cases where more than one significant emission peak was present in the CO spectrum, we have used archival maser or high-density gas tracers to determine the source velocity (see \citealt{urquhart_13co_north} for more details).
In total. we have been able to assign a unique velocity to 300 MYSO candidates and compact HIT regions within the GRS region.
In total, we have been able to assign a unique velocity to 300 MYSO candidates and compact HII regions within the GRS region.
By comparing the Galactic longitudes and latitudes and the velocities of the RMS sources with those derived for the clouds reported by ?. we are able to identify the population of clouds that are responsible for giving birth to the next generation of massive stars in the Galaxy.
By comparing the Galactic longitudes and latitudes and the velocities of the RMS sources with those derived for the clouds reported by \citet{rathborne2009} we are able to identify the population of clouds that are responsible for giving birth to the next generation of massive stars in the Galaxy.
To find a match we began by searching within the GRS in a small region around the coordinates of each RMS source.
To find a match we began by searching within the GRS in a small region around the coordinates of each RMS source.
This region was 5x5x11 resolution elements in the /xbxv directions corresponding to 15 aremin in Galactic longitude and latitude and 4.67 iin velocity.
This region was $5\times 5 \times 11$ resolution elements in the $l \times b \times v$ directions corresponding to 15 arcmin in Galactic longitude and latitude and 4.67 in velocity.
This search returned a unique cloud match for 116 RMS sources.
This search returned a unique cloud match for 116 RMS sources.
For a further 175 sources. multiple clouds were found within the search radius.
For a further 175 sources, multiple clouds were found within the search radius.
In the majority of these case there were only a few pixels from a neighbouring cloud and. to determine a match. we selected the cloud that contributed the most pixels within the search region.
In the majority of these case there were only a few pixels from a neighbouring cloud and, to determine a match, we selected the cloud that contributed the most pixels within the search region.
We failed to find any cloud association for fifteen of the 306
We failed to find any cloud association for fifteen of the 306
The Xaw tranusicut SAN JLs0s.l3658 was discovered in September 1996 when it exhibited a weak. outbux‘st lasting ouly a few weeks (IutZaudetal.1998. 2001).
The X-ray transient SAX J1808.4–3658 was discovered in September 1996 when it exhibited a weak outburst lasting only a few weeks \citep{intzandetal1998,intzandetal2001}.
. In April 1998 the source was fouud to be in outburst again (Marshall1998) and it was discovered that the source exhibits coherent iuilliseconcd X-rav oscillatious with a frequency ο: approximately 101 Iz (Wijnands&vanderWhis 1998).
In April 1998 the source was found to be in outburst again \citep{marshall1998} and it was discovered that the source exhibits coherent millisecond X-ray oscillations with a frequency of approximately 401 Hz \citep{wvdk1998}.
Iu early 2000. the source exhibited a hird outburst during which it showed erratic Iuniirosity behavior with Iuuinosity swings of three oxers of magnitude within a few davs (Wijnaucdseal.2001.2002)..
In early 2000, the source exhibited a third outburst during which it showed erratic luminosity behavior with luminosity swings of three orders of magnitude within a few days \citep{wijnandsetal2001_rxte,wijnandsetal2002_bepposax}.
This erratic behavior lasted for several mouths before the source reurnued to quiescence.
This erratic behavior lasted for several months before the source returned to quiescence.
Very receuth. in October 2002. a fourth οἱthirst of the source was detected (Markwardt./Miller.Wijnands2002)/— during which its peak Iuminositv was very simular to that observed caving the 1996 aud 1998 oubursts.
Very recently, in October 2002, a fourth outburst of the source was detected \citep{markwardtetal2002} during which its peak luminosity was very similar to that observed during the 1996 and 1998 outbursts.
Iun quiescence, SAN. JLsos.lb3658 has been observed on several occasions with theDeppoSAX and satellites (Stellaetal.2010:Dotaui.Asai.&Wijnaucds2000:Wiinandsetal. 2002)..
In quiescence, SAX J1808.4–3658 has been observed on several occasions with the and satellites \citep{stellaetal2000,daw2000,wijnandsetal2002_bepposax}.
The source was very din iu quiex(nee. with a lDunimositv close to or lower ha1 107st.
The source was very dim in quiescence, with a luminosity close to or lower than $10^{32}$.
Due to the limited angular resolution ofDeppoSANX. doubts were raised as o whether the source detected by this satellite was truly SAN JisUs8.L3658 or an unrelated field source (Wijnaudsοal.2002)..
Due to the limited angular resolution of, doubts were raised as to whether the source detected by this satellite was truly SAX J1808.4–3658 or an unrelated field source \citep{wijnandsetal2002_bepposax}.
Campanactal.(2002) reported on a quiesceut observation o the source performed with which resolved this issue.
\citet{campanaetal2002} reported on a quiescent observation of the source performed with which resolved this issue.
They detected the source at a lininositv of 5s10% aud found that the feld around SAN JLS0s.3658 is rather crowded with weak sources.
They detected the source at a luminosity of $5\times10^{31}$ and found that the field around SAX J1808.4–3658 is rather crowded with weak sources.
Two such sources are relatively close to SANJInüs.3658 and mieht have conceivably caused a systematic positional offset during the
Two such sources are relatively close to SAXJ1808.4–3658 and might have conceivably caused a systematic positional offset during the
(z.5SFH) space.
$z$ $SFR$ ) space.
To be precise. one needs to include a correction term accounting for the cosmic microwave background (CAIB) because most SED cata are measured in contrast to the CAIB.
To be precise, one needs to include a correction term accounting for the cosmic microwave background (CMB) because most SED data are measured in contrast to the CMB.
In practice this correction can be safely ignored as long as 7)/(1+2)2»2.7 kx. As a galaxy is placed further and further away. its entire SED shilis to the bottom (fainter. because of D>) and to the left (lower freeuenev. van.=v/(12-2)) in the (2.5FR) space.
In practice this correction can be safely ignored as long as $T_d/(1+z) \gg 2.7$ K. As a galaxy is placed further and further away, its entire SED shifts to the bottom (fainter, because of $D_L^{-2}$ ) and to the left (lower frequency, $\nu_{obs}=\nu_o/(1+z)$ ) in the $z$ $SFR$ ) space.
This generic behavior for a dusty starburst and the resulting change in the apparent spectral index between 1.4 Giz and 350 yam has been pointed out. previously as a redshift inclicator by(1999).
This generic behavior for a dusty starburst and the resulting change in the apparent spectral index between 1.4 GHz and 850 $\mu$ m has been pointed out previously as a redshift indicator by.
. The slope of the rising part of the cust spectrum is such that the Doppler shift of the spectrum nearly offsets the D,? drop in flux density. making the subimin bands particularly attractive for blind searches. but the SED measurements on this part of the dust spectrum offer limited redshift information for (he same reason2001).
The slope of the rising part of the dust spectrum is such that the Doppler shift of the spectrum nearly offsets the $D_L^{-2}$ drop in flux density, making the submm bands particularly attractive for blind searches, but the SED measurements on this part of the dust spectrum offer limited redshift information for the same reason.
. As in other photometric redshift techniques. the redshift inlormation comes from distinct spectral features such as the (rough between the declining non-thermal svnchrotron emission in radio and the sharp rise in the dust spectrum or the dust peak near the rest svavelength of LOO jam. Even when the dust peak in the FIR is not sampled by observations. the radio svnchrotron measurements help set the vertical scale with respect to Che dust spectrum. i.e. the SFR.
As in other photometric redshift techniques, the redshift information comes from distinct spectral features such as the trough between the declining non-thermal synchrotron emission in radio and the sharp rise in the dust spectrum or the dust peak near the rest wavelength of 100 $\mu$ m. Even when the dust peak in the FIR is not sampled by observations, the radio synchrotron measurements help set the vertical scale with respect to the dust spectrum, i.e. the $SFR$.
Because SED measurements for most submm galaxies include only a few discrete points rather (han a continuous lrequency coverage. our best fii SED model search utilizes a v minimization with discrete sampling of the parameter space rather than a full technique.
Because SED measurements for most submm galaxies include only a few discrete points rather than a continuous frequency coverage, our best fit SED model search utilizes a $\chi^2$ minimization with discrete sampling of the parameter space rather than a full cross-correlation technique.
One acvantage of this approach is that the upper limits in fIux density can be incorporated in a straightforward wav by simply rejecting all (ial SEDs that are incompatible with the upper limits.
One advantage of this approach is that the upper limits in flux density can be incorporated in a straightforward way by simply rejecting all trial SEDs that are incompatible with the upper limits.
To test the robustness of our SED template. we apply this photometric redshift technique to several well studied submm galaxies. and these results are summarized in Table 2..
To test the robustness of our SED template, we apply this photometric redshift technique to several well studied submm galaxies, and these results are summarized in Table \ref{tab:ztable}.
Disregarding the formal uncertainties lor the moment. the new photometric redshifts z,j, are in excellent agreement with the spectroscopic redshifts τιν for the all live «ΠΡΙ galaxies wilh known redshifts.
Disregarding the formal uncertainties for the moment, the new photometric redshifts $z_{ph}$ are in excellent agreement with the spectroscopic redshifts $z_{sp}$ for the all five submm galaxies with known redshifts.
When compared to the old radio-to-subnun spectral index estimates zsy as shown in Fieure 4.. the improvement is seen mainiv at high redshift (2> 2) where the effectiveness of the spectral index method diminishes due to the flattening of the a—z relation.
When compared to the old radio-to-submm spectral index estimates $z_{SI}$ as shown in Figure \ref{fig:comparez}, the improvement is seen mainly at high redshift $z>2$ ) where the effectiveness of the spectral index method diminishes due to the flattening of the $\alpha-z$ relation.
Some improvement is generally expected for the new photometric method. since more information is utilized.
Some improvement is generally expected for the new photometric method since more information is utilized.
At the same time this comparison also highlights the efficiency of the spectral index technique which utilizes rather limited amount of information.
At the same time this comparison also highlights the efficiency of the spectral index technique which utilizes rather limited amount of information.
is worth of mentioning (hat. since we are performing a strictly differenal analvsis between the SGD-a and the dominant cluster population. small uncertainties in (he reddening; aud distance modulus would not affect the overall results.
is worth of mentioning that, since we are performing a strictly differential analysis between the SGB-a and the dominant cluster population, small uncertainties in the reddening and distance modulus would not affect the overall results.
As can be appreciated in Fig. 6..
As can be appreciated in Fig. \ref{fig:fit1},
while the metal-poor population is reasonably well reproduced. the TO level of (he anomalous population is definitely fainter (han what predicted by a metal-rich isochrone significantly vounger (/«146yr) than the dominant population.
while the metal-poor population is reasonably well reproduced, the TO level of the anomalous population is definitely fainter than what predicted by a metal-rich isochrone significantly younger $t<14 Gyr$ ) than the dominant population.
A detailed analvsis of Fig.
A detailed analysis of Fig.
6 shows that indeed the observed morphology of the anomalous SGB is not correctly reproduced by anv Z=0.005 isochrone. regardless of age.
\ref{fig:fit1} shows that indeed the observed morphology of the anomalous SGB is not correctly reproduced by any $Z=0.005$ isochrone, regardless of age.
In. fact. though the MS-TO level of the anomalous population appears consistent with an age of 17 Gyr. (Le. Ofe2 Gyr older than the metal-poor component). the TO color is significantly bluer (han the corresponding isochrone.
In fact, though the MS-TO level of the anomalous population appears consistent with an age of 17 Gyr, (i.e. $\delta t \sim 2$ Gyr older than the metal-poor component), the TO color is significantly bluer than the corresponding isochrone.
Moreover. (he morphology. extension aud position of the lower RGB-a and SGD-a are nol correctly reproduced. since the SGB-a appears significantlv less extended in color and much steeper than what predicted by a metal rich isochrone (somewhat suggestive of a lower metal content): the position of the base of the RGB-a appears significantly bluer than what predicted by the The same problems have been noted by fitting the SGB-a in the (V.D-V) plane.
Moreover, the morphology, extension and position of the lower RGB-a and SGB-a are not correctly reproduced, since the SGB-a appears significantly less extended in color and much steeper than what predicted by a metal rich isochrone (somewhat suggestive of a lower metal content); the position of the base of the RGB-a appears significantly bluer than what predicted by the The same problems have been noted by fitting the SGB-a in the (V,B-V) plane.
Note that the same isochrone set nicely fits (he SGD/TO region of well-known clusters of similar metallicity (as 47 Tuc).
Note that the same isochrone set nicely fits the SGB/TO region of well-known clusters of similar metallicity (as 47 Tuc).
Other (wo parameters. namely (he He and CNO abuuelances can alfect the location and shape of the SGB-TO region in the CMD.
Other two parameters, namely the He and CNO abundances can affect the location and shape of the SGB-TO region in the CMD.
In order to evaluate the effect of an enhanced 116 abundance over the shape of the SGD. we computed a set of suitable models al Y=0.28 (0Y—0.05 with respect to the standard value adopted above).
In order to evaluate the effect of an enhanced He abundance over the shape of the SGB, we computed a set of suitable models at Y=0.28 $\delta Y=0.05$ with respect to the standard value adopted above).
From this models we noted that an inerease of the He abundance does not significantly affect the TO level while it decreases the SGD extension by moving blueward the location of the RGB by few hundredths ol magnitude in the 7—R color.
From this models we noted that an increase of the He abundance does not significantly affect the TO level while it decreases the SGB extension by moving blueward the location of the RGB by few hundredths of magnitude in the $B-R$ color.
Hence an increase of the He abundance cannot be invoked to explain the faintness of the TO of the anomalous population.
Hence an increase of the He abundance cannot be invoked to explain the faintness of the TO of the anomalous population.
Renzini (1977. see also the discussion in Salaris. Chiefli Stvamero 1993) showed that the TO Iumninosity and temperature primarily depend on the CNO and Ne abundances. while the RGB colors mainly depend on the Fe. $i ancl Mg ones.
Renzini (1977, see also the discussion in Salaris, Chieffi Straniero 1993) showed that the TO luminosity and temperature primarily depend on the CNO and Ne abundances, while the RGB colors mainly depend on the Fe, Si and Mg ones.
Thus anhoc increase of Zeo,v« would. move the TO color towards the red while leaving (he RGB location unchanged. thus reducing the SGB extension ancl move the TO towards fainter magnitudes (see Fig.
Thus an increase of $Z_{C,N,O,Ne}$ would move the TO color towards the red while leaving the RGB location unchanged, thus reducing the SGB extension and move the TO towards fainter magnitudes (see Fig.
22 of Salaris. Clhieffi Straniero 1993).
22 of Salaris, Chieffi Straniero 1993).
This
This
2001).
.
. Were we describe a further important extension to verv low mass stars. all of which are within the ONC defined on ονποσα grounds by Hillenbrand&Iartmann(1993).
Here we describe a further important extension to very low mass stars, all of which are within the ONC defined on dynamical grounds by \citet{hh98}.
.. This is (hie relevant. portion of the Orion association [or direct comparison with older open clusters since it is Cae only part (hat will likely maintain its identity for 50 My or longer.
This is the relevant portion of the Orion association for direct comparison with older open clusters since it is the only part that will likely maintain its identity for 50 My or longer.
Also. it is important (o isolate a sample with as small an age range as possible in (hese studies. since rotation periods may evolve rapidly during the first ~1-2 My of a star's life.
Also, it is important to isolate a sample with as small an age range as possible in these studies, since rotation periods may evolve rapidly during the first $\sim$ 1-2 My of a star's life.
For these reasons it is best to focus on the ONC. as opposed to the entire Orion A population (or other T associations). when studying the time dependence of stellar angular momentum using clusters.
For these reasons it is best to focus on the ONC, as opposed to the entire Orion A population (or other T associations), when studying the time dependence of stellar angular momentum using clusters.
Ninety-two images of the ONC were obtained through an intermediate band. filler (A, =815.9 nm: AAyyy=20.9 nm: selected to exelude strong nebular lines) on 45 nights between 25 Dec. 1998 and 22 Feb. 1999 with the Wide Field Imager attached to the 2.2m MPG/ESO telescope at La Silla in Chile.
Ninety-two images of the ONC were obtained through an intermediate band filter $\lambda_c$ =815.9 nm; $\Delta \lambda_{FWHM}$ =20.9 nm; selected to exclude strong nebular lines) on 45 nights between 25 Dec. 1998 and 22 Feb. 1999 with the Wide Field Imager attached to the 2.2m MPG/ESO telescope at La Silla in Chile.
Details of the dala acquisition. analvsis ancl additional results will be reported elsewhere (IIerbstοἱal.2001).
Details of the data acquisition, analysis and additional results will be reported elsewhere \citep{her01}.
. Here we note that the [ield surveved was a 33! x 34 rectangle centered approximately on 6! Ori C. making it nearly coincident with the ONC clefined by Hillenbrand. and Hartmann (1998).
Here we note that the field surveyed was a $33\arcmin$ x $34\arcmin$ rectangle centered approximately on $\theta^1$ Ori C, making it nearly coincident with the ONC defined by Hillenbrand and Hartmann (1998).
Photometry was obtained on 2.294 stars extending to 19918.
Photometry was obtained on 2,294 stars extending to $\sim$ 18.
A search for periodicity was carried out. using the Lomb-Scargle technique and standard assessments of false alarm probabilities elal.2000:Rebull 2001).
A search for periodicity was carried out using the Lomb-Scargle technique and standard assessments of false alarm probabilities \citep{her00,reb01}.
. Of the 404 periodic stars identified in our sample. 111 had been previously discovered bv Herbstetal.(2000) or Stassunetal.(1999). and 99 of thesewere found to have identical periods to within the errors of the determinations (οὐ ο).
Of the 404 periodic stars identified in our sample, 111 had been previously discovered by \citet{her00} or \citet{sta99} and 99 of thesewere found to have identical periods to within the errors of the determinations $\sim$ ).
Α list of periodic stars is available upon request (ο the first author.
A list of periodic stars is available upon request to the first author.
Eleven of the 12 stars which had. period disagreements were cases of either harmonics or beat periods masquerading as fundamentals.
Eleven of the 12 stars which had period disagreements were cases of either harmonics or beat periods masquerading as fundamentals.
The quality aud quantity of the data obtained lor (his study is unprecedented and (hat is responsible for our success in nearly (ripling (he number of rotation periods now known for stars in (he ONC.
The quality and quantity of the data obtained for this study is unprecedented and that is responsible for our success in nearly tripling the number of rotation periods now known for stars in the ONC.
In particular. our study extends to fainter stars and. therefore. lower masses (han have been probed heretofore ancl it is that aspect of the results which we primarily discuss in thisLetter.
In particular, our study extends to fainter stars and, therefore, lower masses than have been probed heretofore and it is that aspect of the results which we primarily discuss in this.
Fie.
Fig.
1 shows the distribution of rotation periods for 335 periodic stars im our sample which have masses determined by Hillenbrand(1997).
1 shows the distribution of rotation periods for 335 periodic stars in our sample which have masses determined by \citet{hil97}.
. Her determinations are based on a comparison of each star's location in the IIR. diagram with PAIS models of (1994).
Her determinations are based on a comparison of each star's location in the HR diagram with PMS models of \citet{dm94}.
. Masses are. of course. model-dependent ancl these could be systematically in error by perhaps as much as 50%.
Masses are, of course, model-dependent and these could be systematically in error by perhaps as much as .
. However. all models of PAIS stars indicate that lower
However, all models of PMS stars indicate that lower
Section 3.2 remove this aud other effects.
Section \ref{sec:corrections} remove this and other effects.
Dark frames aud snuadl scale flat field features were stable over a few cays to a week.
Dark frames and small scale flat field features were stable over a few days to a week.
A few bad columns in the CCDs did not affect the overall quality of the data set.
A few bad columns in the CCDs did not affect the overall quality of the data set.
Cameras did not perform equally.
Cameras did not perform equally.
The sliehtlv. lower photometric quality of camera can be seen in survey statistics presented i Section L1..
The slightly lower photometric quality of camera can be seen in survey statistics presented in Section \ref{sec:survey}.
The loss of observing time due to temporary failure of cauucera is also visible.
The loss of observing time due to temporary failure of camera is also visible.
Despite the fact that the primary goal of the ROTSE-I project was rapid response to GRB trigeers aud not sky patrols. alinost all observing time was actually speut iu the latter mode.
Despite the fact that the primary goal of the ROTSE-I project was rapid response to GRB triggers and not sky patrols, almost all observing time was actually spent in the latter mode.
An accessible CRB position would be posted bv the GCN network approximately once every 10 dave.
An accessible GRB position would be posted by the GCN network approximately once every 10 days.
Upon receipt of the coordinates. the ROTSE-I system would abort the curent patrol activity and iunmediately start observing the field around tle position or approximately one hour of imagine.
Upon receipt of the coordinates, the ROTSE-I system would abort the current patrol activity and immediately start observing the field around the position for approximately one hour of imaging.
At the beeinnine of each nieht. about 12 dark frames were collected for calibration purposes.
At the beginning of each night, about 12 dark frames were collected for calibration purposes.
No special flat Ποια exposures were made (Section 3.1.1)).
No special flat field exposures were made (Section \ref{sec:basic}) ).
The larec combined field of view delivered ly the ROTSE-I svstem requires only 206 tiles to cover the cutive sky. with 161 iles observable from Los Alamos.
The large combined field of view delivered by the ROTSE-I system requires only 206 tiles to cover the entire sky, with 161 tiles observable from Los Alamos.