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Also shown in Figure Lo is the imass-depeudoeut age at which a star transitions from the C sequence to the I sequence. | Also shown in Figure \ref{regimes2} is the mass-dependent age at which a star transitions from the C sequence to the I sequence. |
The fox of these two. indepeudently-determined. empirical transitions is remarkably simular. | The form of these two, independently-determined, empirical transitions is remarkably similar. |
This iuplics that stars with saturated X-ray ciissiou are found on the € sequence of fast rotators. while those with non-saturated cinission can be found on the I sequence of slow rotators. | This implies that stars with saturated X-ray emission are found on the C sequence of fast rotators, while those with non-saturated emission can be found on the I sequence of slow rotators. |
A similar result to this was already hinted at frou, a πιο. smaller sample bv ?) who associated the presence of stars on the rotational I sequence with unsaturated X-ray cutters. but suggested that stars on the € sequence were coronallyv super-saturated and stars in the eap between the C and I sequences showed saturated X-ray. emission. | A similar result to this was already hinted at from a much smaller sample by \citet{barn03b} who associated the presence of stars on the rotational I sequence with unsaturated X-ray emitters, but suggested that stars on the C sequence were coronally super-saturated and stars in the gap between the C and I sequences showed saturated X-ray emission. |
The association of supersaturated coronal emitters with stars ou the C sequence can be argued against simply based ou the very small ΠΡΟ of super-saturated A-rav cutters (see Section 3.3)) compared to the unich larger uuuber of fast rotators found on the € sequence of vouug clusters. | The association of super-saturated coronal emitters with stars on the C sequence can be argued against simply based on the very small number of super-saturated X-ray emitters (see Section \ref{s-supersaturation}) ) compared to the much larger number of fast rotators found on the C sequence of young clusters. |
The correlation between the trausitiou from saturated to uusaturated coronal enission aud from the rotational C sequence to the I sequence is not surprising when oue considers that both angular 1iomienutuim loss aud coronal A-rav chussion are products of the stellar magnetic dynamo aud known to scale with the convective turnover nie (eg.7). | The correlation between the transition from saturated to unsaturated coronal emission and from the rotational C sequence to the I sequence is not surprising when one considers that both angular momentum loss and coronal X-ray emission are products of the stellar magnetic dynamo and known to scale with the convective turnover time \citep[e.g.][]{mont01}. |
This scaling was noted by 7?) who showed that the functional dependence of the angular nomenti loss rate was related to the convective "urmnover time. and henceforth derived a new model for he rotational evolution of cool stars in terms of the Rossby uiuuber (2).. | This scaling was noted by \citet{barn10a} who showed that the functional dependence of the angular momentum loss rate was related to the convective turnover time, and henceforth derived a new model for the rotational evolution of cool stars in terms of the Rossby number \citep{barn10b}. |
Iu fact the balance between coronal osses and angular momentum loss is believed to be due o the configuration of maeuetic field lines at the stellar surface. with the fraction of open magnetic field. lines deteriuuius the spin-down rate. while closed maguetic field lines dictate the N-vav huninosity (6.8.?).. | In fact the balance between coronal losses and angular momentum loss is believed to be due to the configuration of magnetic field lines at the stellar surface, with the fraction of open magnetic field lines determining the spin-down rate, while closed magnetic field lines dictate the X-ray luminosity \citep[e.g.][]{holz07}. |
The similarity between two indepeudently determined Clupizical transition criteria is unlikely to be a coincidence. | The similarity between two independently determined empirical transition criteria is unlikely to be a coincidence. |
It therefore sugeests that the changes that occur Within a star as it transitions from the € sequence to the I sequence are also responsible for the star leaving the saturated regine of N-rav cinission. | It therefore suggests that the changes that occur within a star as it transitions from the C sequence to the I sequence are also responsible for the star leaving the saturated regime of X-ray emission. |
One could also consider the opposite causality. that the mechanisis responsible for the changes in X-ray enission also cause changes in the augular momenutiui loss rate. | One could also consider the opposite causality, that the mechanisms responsible for the changes in X-ray emission also cause changes in the angular momentum loss rate. |
IToswever. the evolution of angular momentum loss is a more complex subject that must explain other observable features (6.8. the sequence of young ultrafast rotators iu clusters) and it therefore seems unlikely that the mechamisim responsible for N-rav saturation ds responsible for these effects as well. | However, the evolution of angular momentum loss is a more complex subject that must explain other observable features (e.g. the sequence of young ultrafast rotators in clusters) and it therefore seems unlikely that the mechanism responsible for X-ray saturation is responsible for these effects as well. |
The hypothesis that we are therefore left with is that the processes governing stellar angular momentum loss are also responsible for N-rav saturation. | The hypothesis that we are therefore left with is that the processes governing stellar angular momentum loss are also responsible for X-ray saturation. |
7) argued that the C and I sequences are due to the —coupling of the stellar wind to. respectively. just the convective zone (Which is decoupled from the radiative zone). aud to the entire star. | \citet{barn03} argued that the C and I sequences are due to the coupling of the stellar wind to, respectively, just the convective zone (which is decoupled from the radiative zone), and to the entire star. |
The transition between these zones (across the ceap in the colorrotation period diaerauu) is therefore associated with the coupling of the radiative and couvective zones to each other. | The transition between these zones (across the `gap' in the color–rotation period diagram) is therefore associated with the coupling of the radiative and convective zones to each other. |
7)| suggests that stars on the convective sequence generate a convective or turbulent dvuamo (e.g.2) aud that the shear between the fast spiunius radiative interior aud the couvective euvelope eventually eeuerates an interface dvuamo that results in the transition onto the I sequence. | \citet{barn03} suggests that stars on the convective sequence generate a convective or turbulent dynamo \citep[e.g.][]{durn93} and that the shear between the fast spinning radiative interior and the convective envelope eventually generates an interface dynamo that results in the transition onto the I sequence. |
?) lave since argued that the form of the rotational isoclirones does not agree with the form expected for the relevant | \citet{barn10a} have since argued that the form of the rotational isochrones does not agree with the form expected for the relevant |
The last decade has seen an avalanche of observations of magnetohydrodynamic (MHD) waves in the solar atmosphere. | The last decade has seen an avalanche of observations of magnetohydrodynamic (MHD) waves in the solar atmosphere. |
It is clear now that MHD waves are ubiquitous in the solar atmosphere. | It is clear now that MHD waves are ubiquitous in the solar atmosphere. |
This has triggered new theoretical research for explaining and interpreting the observed properties. | This has triggered new theoretical research for explaining and interpreting the observed properties. |
A special point of attention is whether these MHD waves are slow. fast or Alfvénn waves. | A special point of attention is whether these MHD waves are slow, fast or Alfvénn waves. |
Apparently. a large fraction of the solar MHD waves community favours very clear cut divisions and does not seem to appreciate the possibility of MHD waves with mixed properties. | Apparently, a large fraction of the solar MHD waves community favours very clear cut divisions and does not seem to appreciate the possibility of MHD waves with mixed properties. |
The transverse oscillations observed in coronal loops (seeforexampleAschwandenetal..1999).. often triggered by a nearby solar flare. are interpreted as fast kink MHD waves. | The transverse oscillations observed in coronal loops \citep[see for example][]{Aschwanden1999}, often triggered by a nearby solar flare, are interpreted as fast kink MHD waves. |
A striking property of these transverse waves is their fast damping with damping times of the order of 3 - 5 periods. | A striking property of these transverse waves is their fast damping with damping times of the order of 3 - 5 periods. |
Resonant absorption is up to today the only damping mechanism that offers a consistent explanation of this rapid damping. | Resonant absorption is up to today the only damping mechanism that offers a consistent explanation of this rapid damping. |
Resonant absorption relies on the transfer of energy from a global MHD wave to local resonant. Alfvénn waves. | Resonant absorption relies on the transfer of energy from a global MHD wave to local resonant Alfvénn waves. |
If this mechanism is indeed operational then. this means that the observed transverse oscillations have Alfvénnic properties in at least part of the oscillating loop. | If this mechanism is indeed operational then this means that the observed transverse oscillations have Alfvénnic properties in at least part of the oscillating loop. |
The debate on the nature of MHD waves in the solar atmosphere has gained new momentum when several groups. e.g. (2007). Okamotoetal. (2007)... Tomezyketal.(2007) reported the detection of Alfvénn waves in HINODE observations. | The debate on the nature of MHD waves in the solar atmosphere has gained new momentum when several groups, e.g. \cite{DePontieu2007}, \cite{Okamoto2007}, \cite{Tomczyk2007} reported the detection of Alfvénn waves in HINODE observations. |
Subsequently VanDoorsselaereetal.(2008) argued strongly against the possible presence of Alfvénn waves in the solar corona emphasising that Alfvénn waves cannot be but torsional. | Subsequently \cite{VanDoorsselaere2008} argued strongly against the possible presence of Alfvénn waves in the solar corona emphasising that Alfvénn waves cannot be but torsional. |
VanDoorsselaereetal.(2008) compared fast kink MHD waves to torsional Alfvénn waves and concluded that the HINODE observations have nothing to do with Alfvénn waves but can be explained in terms of fast kink MHD waves. | \cite{VanDoorsselaere2008} compared fast kink MHD waves to torsional Alfvénn waves and concluded that the HINODE observations have nothing to do with Alfvénn waves but can be explained in terms of fast kink MHD waves. |
This paper will not try to explain the HINODE observations. | This paper will not try to explain the HINODE observations. |
Its aim is to determine the nature of kink MHD waves on magnetic flux tubes. | Its aim is to determine the nature of kink MHD waves on magnetic flux tubes. |
We have no doubt about the explanation of transverse oscillation of coronal loops in terms of kink MHD waves. | We have no doubt about the explanation of transverse oscillation of coronal loops in terms of kink MHD waves. |
MHD waves with their azimuthal wave number equal to I. i.e. 74=I. are the only motions that displace the axis of the loop and the loop as a whole. | MHD waves with their azimuthal wave number equal to 1, i.e. $m=1$, are the only motions that displace the axis of the loop and the loop as a whole. |
It is not clear to us OI what arguments the use of the adjective fast i$ based. | It is not clear to us on what arguments the use of the adjective fast is based. |
As far as we know there has not been any study of the forces that drive the kink waves in coronal loops. | As far as we know there has not been any study of the forces that drive the kink waves in coronal loops. |
If these waves are fast. then the pressure gradient force should be. in general. the dominant force compared to the magnetic tension force. | If these waves are fast, then the pressure gradient force should be, in general, the dominant force compared to the magnetic tension force. |
We have to admit that we also have used the adjective fast without a solid argument in favour of this classification. | We have to admit that we also have used the adjective fast without a solid argument in favour of this classification. |
Our aim is to understand the spatial structure of the motions in the kink waves. | Our aim is to understand the spatial structure of the motions in the kink waves. |
An MHD wave on an axi-symmetric 1-D. cylindrical plasma equilibrium ts characterised by two wave numbers. the azimuthal wave number 77. and the axial wave number K.. | An MHD wave on an axi-symmetric 1-D cylindrical plasma equilibrium is characterised by two wave numbers, the azimuthal wave number $m$, and the axial wave number $k_z$. |
In addition modes can have a different number of nodes in the radial direction and this number can be used to further classify the modes. | In addition modes can have a different number of nodes in the radial direction and this number can be used to further classify the modes. |
Hence. an MHD eigenmode is characterised by three numbers. | Hence, an MHD eigenmode is characterised by three numbers. |
The azimuthal wave number ts an integer. | The azimuthal wave number is an integer. |
The modes with a=0 are usually called sausage (slow and fast) or torsional (Alfvénn). | The modes with $m=0$ are usually called sausage (slow and fast) or torsional (Alfvénn). |
The modes with i=| are named kink and the modes with w>2 are flute modes. | The modes with $m=1$ are named kink and the modes with $m \geq 2$ are flute modes. |
The axial wave number k. can be diseretised as κ.=nF with L the length of the loop and n=1.2.... Depending on the dimensions of the equilibrium model there can be more than one radial eigenmode for a given couple (71.&-). | The axial wave number $k_z$ can be discretised as $k_z = n \frac{\pi}{L}$ with $L$ the length of the loop and $n
=1,2, \ldots$ Depending on the dimensions of the equilibrium model there can be more than one radial eigenmode for a given couple $(m, k_z)$. |
In what follows we shall study linear MHD waves that are superimposed on a flux tube in static equilibrium with a straight and constant axial magnetic field. | In what follows we shall study linear MHD waves that are superimposed on a flux tube in static equilibrium with a straight and constant axial magnetic field. |
This equilibrium model contains the essential physics of the problem and allows a relatively straightforward mathematical analysis of the MHD waves. | This equilibrium model contains the essential physics of the problem and allows a relatively straightforward mathematical analysis of the MHD waves. |
MHD waves have been investigated in previous studies. | MHD waves have been investigated in previous studies. |
However. these studies almost exclusively focused on the frequencies of the MHD waves and in addition they were in most cases restricted to real frequencies. | However, these studies almost exclusively focused on the frequencies of the MHD waves and in addition they were in most cases restricted to real frequencies. |
For example. the paper by Edwin&Roberts (1983).. which is often referred to in the solar MHD wave community. is limited to real frequencies and does not give any information on the eigenfunctions beyond the fact that they can be expressed in terms of Bessel functions. | For example, the paper by \cite{Edwin1983}, which is often referred to in the solar MHD wave community, is limited to real frequencies and does not give any information on the eigenfunctions beyond the fact that they can be expressed in terms of Bessel functions. |
Complex frequencies were considered by Spruit(1982) and by | Complex frequencies were considered by \cite{Spruit1982} and by |
observational frames. | observational frames. |
The secoud aperture correction Cs is au ollse of -0.10 (irrespective of filter and detector) to convert the magnitude [rom the 075 radius into a iomuinal infinite aperture. | The second aperture correction $C_{\infty}$ is an offset of –0.10 (irrespective of filter and detector) to convert the magnitude from the $\farcs$ 5 radius into a nominal infinite aperture. |
The zero poiuts ZPy are taken [roin Baggett et al. ( | The zero points $ZP_{\rm V}$ are taken from Baggett et al. ( |
1997). | 1997). |
We have also considered the efficiency (CTE) correction Ceyp in its new lormulation » Whitino'e. Hever. Casertano (1990). | We have also considered the (CTE) correction $C_{CTE}$ in its new formulation by Whitmore, Heyer, Casertano (1999). |
Other secoudary calibraious have been neglected. | Other secondary calibrations have been neglected. |
The correction with the new rates derived (rom Baggett Gouzaga (1998) is practically zeὉ for all tle detectors in the two E555W aud FSIIW filters. | The correction with the new rates derived from Baggett Gonzaga (1998) is practically zero for all the detectors in the two F555W and F814W filters. |
We lave justeacl not applied auy correction Lor heanomaly (see e.g. Casertano Mutcller 1998). because its value has recently tu‘ned out to be more uncertain than previously thought (Casertano. 2000 private comamuunicatiou). | We have instead not applied any correction for the (see e.g. Casertano Mutchler 1998), because its value has recently turned out to be more uncertain than previously thought (Casertano, 2000 private communication). |
For a sale interpreation of the characteristics of tle galaxy stellar populations. we ueed to distinguish as nach as possible sigle stars [rom exteuded. bleuded or spurious objects. | For a safe interpretation of the characteristics of the galaxy stellar populations, we need to distinguish as much as possible single stars from extended, blended or spurious objects. |
We have thus appliec to our catalogs selection criteria based ou the shape of the objects. | We have thus applied to our catalogs selection criteria based on the shape of the objects. |
To this aim we have considered. the Daophot parameters V7> audsharpness: V77 gives. the ratio. of+ the observed scater iu the fit residuals to the expected scatter calculated from a predictive model based ou the measurect detector features. whilesherpness sets the iirisic angular size of the objects. | To this aim we have considered the Daophot parameters $\chi^2$ and: $\chi^2$ gives the ratio of the observed pixel-to-pixel scatter in the fit residuals to the expected scatter calculated from a predictive model based on the measured detector features, while sets the intrinsic angular size of the objects. |
We lave seected only the objects with 4?€ Lin both filters for the PC. 4?<2.2 in F555W and FS1IW for the three WEs. | We have selected only the objects with $\chi^2\leq\,$ 4 in both filters for the PC, $\chi^2\leq\,$ 2.2 in F555W and $\chi^2\leq$ 2.5 in F814W for the three WFs. |
Moreover only detections with €sharpnessx 0.5 iu all filters aid for all detectors have been eventually retained. | Moreover only detections with $\,\leq\,sharpness\,\leq\,$ 0.5 in all filters and for all detectors have been eventually retained. |
These X? andsharpness values turned out to be tlose allowing to best reject spurious and extended objects. without eliminating also the wieght stars. | These $\chi^2$ and values turned out to be those allowing to best reject spurious and extended objects, without eliminating also the bright stars. |
By iuspecting the rejected objects. we have recognized several caudidate star clusters (i.e. fairly round but extended ojets) aud background galaxies. | By inspecting the rejected objects, we have recognized several candidate star clusters (i.e. fairly round but extended objects) and background galaxies. |
We have 17 star clusters (besides he brightest SSC) in the PC. fin the ΛΕΣ. 5 in the WE3 and 3 in the WEf. | We have 17 star clusters (besides the brightest SSC) in the PC, 4 in the WF2, 5 in the WF3 and 3 in the WF4. |
The candidate ockgrouud galaxies are 13 in the WE2. 11 in the ΝΕΟ. and 1s in the WFf. | The candidate background galaxies are 13 in the WF2, 11 in the WF3, and 18 in the WF4. |
In the latter case. the 18 galaxies occupy only a restricted zone of the WEI auc the conceutratiou is much higher tliau in the other fields: we ae presumably dealing with a real. vet uuclassilied. galaxy cluster. since the derived galaxy density (15 ~ down to 7222) is almost a factor of 5 higher than that oredicted by deep couns of field galaxies (e.g. Pozzetti et al. | In the latter case, the 18 galaxies occupy only a restricted zone of the WF4 and the concentration is much higher than in the other fields: we are presumably dealing with a real, yet unclassified, galaxy cluster, since the derived galaxy density (18 $^{-2}$ down to 22) is almost a factor of 5 higher than that predicted by deep counts of field galaxies (e.g. Pozzetti et al. |
1998). | 1998). |
A further selection olten applied is based on the photonetric error op4o. | A further selection often applied is based on the photometric error $\sigma_{DAO}$. |
The distributior of 7pjo in the F5522W a[i| ESLIW. bauds is shown in Fig. | The distribution of $\sigma_{DAO}$ in the F555W and F814W bands is shown in Fig. |
16 for both the PC aud the three WEs. | \ref{dao_sig} for both the PC and the three WFs. |
[If we apply the opie selection criterion o the stars whicl have already been constrained iu the u and parajeters. witl συιοx 0.2 in both [ilte ‘swe retain 15299 stars in the PC. and 67825. 2812 and 1279 i the WE2. ΛΕΡ aia4 WFL respectively. | If we apply the $\sigma_{DAO}$ selection criterion to the stars which have already been constrained in the $\chi^2$ and parameters, with $\sigma_{DAO}\leq\,$ 0.2 in both filters we retain 15299 stars in the PC, and 6785, 2812 and 4279 in the WF2, WF3 and WF4, respectively. |
Going down to συ4o <Q.1 implies 922| stars retained iu he PC. and L691. 1551 aud 2376 in the WE2. ΛΕΡ aud WFfL respectively. | Going down to $\sigma_{DAO}\leq\,$ 0.1 implies 9224 stars retained in the PC, and 4691, 1551 and 2376 in the WF2, WF3 and WF4, respectively. |
We have put all te measured stars οἱ the common reference frame shown in tle mosale image of Fig. 1L. | We have put all the measured stars on the common reference frame shown in the mosaic image of Fig. \ref{mosaico_bw}, |
ard we have divided them iutO six groups according t) their position with respect to the isophotal «'ontour evels. | and we have divided them into six groups according to their position with respect to the isophotal contour levels. |
Regious 1 an 2, aud Regions f aud 5. are considered together because both their stellar poplatious aid their crowcdiug conditions turned out to be very similar to each other. | Regions 1 and 2, and Regions 4 and 5, are considered together because both their stellar populations and their crowding conditions turned out to be very similar to each other. |
Region 7. despite the low uumber of stars. has beencousidered separately from the others. because it contains tle brightest SSC aud crowding dramatically allects its photometric accuracy | Region 7, despite the low number of stars, has beenconsidered separately from the others, because it contains the brightest SSC and crowding dramatically affects its photometric accuracy |
lines tracing high-mass star formation including 1120 masers. HC4N and CLIOLL. | lines tracing high-mass star formation including $_2$ O masers, $_3$ N and $_3$ OH. |
Observations were carried out on the Mopra racio telescope in Mas/June of 2008 and April of 2009 emploving the Mopra spectrometer (MODS) in zoom-moce. | Observations were carried out on the Mopra radio telescope in May/June of 2008 and April of 2009 employing the Mopra spectrometer (MOPS) in zoom-mode. |
Mopra ids a 22mm single-dish radio telescope (31716'04"S. οοσο. S66m z.s.L.) | Mopra is a m single-dish radio telescope $31^\circ 16^\prime 04^{\prime\prime}$ S, $149^\circ 05^\prime 59^{\prime\prime}$ E, 866m a.s.l.) |
located 50 κι north west of Sydney. Australia. | located $\sim$ km north west of Sydney, Australia. |
The mmm receiver operating in the frequency range of (11112. coupled with the UNSW Alopra wide-bandwidth spectrometer (MODPS). allows an instantaneous GGlIz bandwidth. | The mm receiver operating in the frequency range of GHz, coupled with the UNSW Mopra wide-bandwidth spectrometer (MOPS), allows an instantaneous GHz bandwidth. |
This gives Mopra the ability to cover most of the mmm band and. simultaneously observe many spectral lines. | This gives Mopra the ability to cover most of the mm band and simultaneously observe many spectral lines. |
The zoom-mode of MOPS allows observations from up to 16 windows simultaneously. where cach window is MMlIz wide and contains 4096 channels in cach of two polarisations. | The zoom-mode of MOPS allows observations from up to 16 windows simultaneously, where each window is MHz wide and contains 4096 channels in each of two polarisations. |
At 12mmm this gives MOPS an elfective. bandwidth of with resolution kms. | At mm this gives MOPS an effective bandwidth of $\sim$ with resolution $\sim$ . |
Within this band the Mopra. beam. FWILAL varies from 2.4 GClIz) to L7' CCH) (Urquhart 2010). | Within this band the Mopra beam FWHM varies from $^\prime$ GHz) to $^\prime$ GHz) \citep{mopra_beam}. |
. Listed in Table 1. are some of the lines which are simultaneously within the bandpass in our configuration. | Listed in Table \ref{tab:lines} are some of the lines which are simultaneously within the bandpass in our configuration. |
On-the-lly Alappine (OTE) observations were conducted in May. of 2008. ancl consisted. of four regions. which are referred. as X. Οοπσ 259). B. (20 20) and C (15. 15) to. cover the Γον emission. peaks from Πο 1500-240 (Aharonianetal.2008b) and map D (20°. 20) to cover the TeV emission from LESS 1801-233 (Aharonianetal.2008b). | On-the-fly Mapping (OTF) observations were conducted in May of 2008, and consisted of four regions, which are referred as A $^{\prime}$ $\times$ $25^{\prime}$), B $^{\prime}$ $\times$ $20^{\prime}$ ) and C $^{\prime}$ $\times$ $15^{\prime}$ ) to cover the TeV emission peaks from HESS J1800-240 \citep{hess_w28} and map D $^{\prime}$ $\times$ $20^{\prime}$ ) to cover the TeV emission from HESS 1801-233 \citep{hess_w28}. |
. We mapped cach region twice. scanning once in right ascension and once in declination in order to reduce noise levels and to eliminate. artificial stripes that can be introduced. when only one scanning direction is used. | We mapped each region twice, scanning once in right ascension and once in declination in order to reduce noise levels and to eliminate artificial stripes that can be introduced when only one scanning direction is used. |
Lt also allows us to check for artefacts that mav occur in one scan. but not the other. | It also allows us to check for artefacts that may occur in one scan, but not the other. |
For reference. the mapped regions are indicated as dashed boxes in Figure 2.. | For reference, the mapped regions are indicated as dashed boxes in Figure \ref{fig:peakpixmapNH311}. |
The OTE mapping parameters we used are similar to hose used in the the LO Southern Galactic Plane Survey (MOPS). a mmm study of the Galactic Plane (Walshοἱal. 2008).. | The OTF mapping parameters we used are similar to those used in the the $_2$ O Southern Galactic Plane Survey (HOPS), a mm study of the Galactic Plane \citep{hops_pilot}. |
Since WHOS also covered the W28 region. we iive also included LOPS data in our mapping analysis. improving our exposure in the mapped areas DB. € D ον a factor of two. | Since HOPS also covered the W28 region, we have also included HOPS data in our mapping analysis, improving our exposure in the mapped areas B, C D by a factor of two. |
Map A extended: beyond the Galactic atituce limit of LOPS (b= 0.57) and so only partial overlap (~25%)) exists. | Map A extended beyond the Galactic latitude limit of HOPS $b=-0.5^\circ$ ) and so only partial overlap $\sim$ ) exists. |
Based on mapping results ancl prior knowledge of he regions under consideration. follow-up single pointing »osition-switched deep spectra were performed in June 2008 and in April of 2009. to provide high sensitivity to vield accurate measurement of the satellite lines which are necessary to determine the gas temiperature and density. | Based on mapping results and prior knowledge of the regions under consideration, follow-up single pointing position-switched deep spectra were performed in June 2008 and in April of 2009, to provide high sensitivity to yield accurate measurement of the satellite lines which are necessary to determine the gas temperature and density. |
As we show shortly. several regions were found to exhibit and higher transitions with satellite lines apparent in the (1.1) spectra. | As we show shortly, several regions were found to exhibit and higher transitions with satellite lines apparent in the (1,1) spectra. |
Regions of bright emission from each map were targeted. in these deep spectra. | Regions of bright emission from each map were targeted in these deep spectra. |
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