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As expected. if we employ. ft.=30 Myr (case 3]. the resulting mass distribution is log-normal. ancl matches the observed mass distribution (Fig. 2)) | As expected, if we employ $t_{\rm dis}^4 = 30$ Myr (case [3]), the resulting mass distribution is log-normal, and matches the observed mass distribution (Fig. \ref{powerlaw.fig}) ) |
very closely. | very closely. |
This is not surprising. since the determination of this short disruption time-scale was based on the assumption of an initial power-law CME with a slope of 2. | This is not surprising, since the determination of this short disruption time-scale was based on the assumption of an initial power-law CMF with a slope of $-2$. [ |
Ehis shows that the assumption of instantaneous disruption. adopted in de Cirijs et al. ( | This shows that the assumption of instantaneous disruption, adopted in de Grijs et al. ( |
2003a.b) hardly allects the determination of the disruption tinc-scales.] | 2003a,b) hardly affects the determination of the disruption time-scales.] |
Lf, instead. we use an initial power-law CALF as before. but now assume that the longer disruption. time-scale predicted by. Baumgardt Alakino's (2003) N-bocly simulations is correct (cases 1.2]. the evolved. cluster mass distribution shows a (broader) peak. shifted. to lower masses by more than one order of magnitude. | If, instead, we use an initial power-law CMF as before, but now assume that the longer disruption time-scale predicted by Baumgardt Makino's (2003) $N$ -body simulations is correct (cases [1,2]), the evolved cluster mass distribution shows a (broader) peak, shifted to lower masses by more than one order of magnitude. |
Moreover. we argued in the previous section that our estimate of the ambient. density is. likely o have been overestimated bv at [east a [actor of 3. | Moreover, we argued in the previous section that our estimate of the ambient density is likely to have been overestimated by at least a factor of 3. |
Since. in an uncdisturbed tidal field of a galaxy with a logarithmic potential. the disruption time-scale depends on the ambient clensity as faiXpOsου, we have also evolved. a power-law witha L5. larger disruption time-scale (case 4] in Fig. 2)): | Since, in an undisturbed tidal field of a galaxy with a logarithmic potential, the disruption time-scale depends on the ambient density as $t_{\rm
dis} \propto \rho_{\rm amb}^{-0.5}$, we have also evolved a power-law with a $1.7\times$ larger disruption time-scale (case [4] in Fig. \ref{powerlaw.fig}) ); |
his factor of 1.7 allows for the uncertainty in fan between ES and 2.5 M. 7. | this factor of 1.7 allows for the uncertainty in $\rho_{\rm amb}$ between 0.8 and 2.5 $_\odot$ $^{-3}$. |
As expected. this longer. probably more realistic disruption time-scale eives rise to à turnover ocated at a cluster mass smaller than that derived in cases JA] and. 2]. thus strengthening the discrepancy between the | As expected, this longer, probably more realistic disruption time-scale gives rise to a turnover located at a cluster mass smaller than that derived in cases [1] and [2], thus strengthening the discrepancy between the |
sinele crystal the vortex passes from one lattice plane to the adjacent lattice plane. | single crystal the vortex passes from one lattice plane to the adjacent lattice plane. |
Hence. we sel the boundary. conditions such as o—0 αἱz=—o€ and ó=22 al z=ox. | Hence, we set the boundary conditions such as $\phi = 0$ at$z=-\infty$ and $\phi = 2\pi$ at $z =\infty$. |
Then. we can solve equations (16)) analvtically and derive the equilibrium configurations. where Figure 3. shows that the vortex can move from one lattice plane to the adjacent laltice plane bv bending itself ancl forming a kink when (he vortex tension is finite. | Then, we can solve equations \ref{eq:stst-int1}) ) analytically and derive the equilibrium configurations, where Figure \ref{kink} shows that the vortex can move from one lattice plane to the adjacent lattice plane by bending itself and forming a kink when the vortex tension is finite. |
The vortex line is displaced substantially from (he pinning sites in (he kink part. while in the other part the vortex line is straight ancl passes (rough nearly all of the pinning sites along the z-axis. | The vortex line is displaced substantially from the pinning sites in the kink part, while in the other part the vortex line is straight and passes through nearly all of the pinning sites along the z-axis. |
/ eiven bv equation (19)) expresses a length of the kink. | $l$ given by equation \ref{eq:kinkl}) ) expresses a length of the kink. |
In (the middle of a kink the vortex line is inclined against the z-axis bv 2.57. | In the middle of a kink the vortex line is inclined against the z-axis by $2.5^{\circ}$. |
Energy is needed to create a kink from a straight. vortex line since (he length of a vortex line increases and a part of it should overcome the pinning potential. | Energy is needed to create a kink from a straight vortex line since the length of a vortex line increases and a part of it should overcome the pinning potential. |
The self-enerev of vortex lines is equal to the kinetic energv of the velocity field induced by them. | The self-energy of vortex lines is equal to the kinetic energy of the velocity field induced by them. |
The increase ol the vortex sell-energv due to bending is while (he energv necessary {ο overcome (he pinning potential is The length of a kink can be determined from minimizing the sum of AL. and AL, (Link&Epstein1991:Ruderman2000). | The increase of the vortex self-energy due to bending is while the energy necessary to overcome the pinning potential is The length of a kink can be determined from minimizing the sum of $\Delta E_s$ and $\Delta E_p$ \citep{lin91, rud00}. |
. Using the solution (18)) together with i= 0. the total energv necessary to form a kink is calculated as | Using the solution \ref{eq:kink}) ) together with $\psi=0$ , the total energy necessary to form a kink is calculated as |
We can assume also that the probability p(À.) is dominated by the shape of (he filler aud write: Asstuning a constant comoving densitv of galaxies inside the narrow-band filter. the nunmber of galaxies is proportional to (he comoving volume V. (fortheexplicit(hecomovingvolume.see?..andreferences therein).. | We can assume also that the probability $p(\lambda_z)$ is dominated by the shape of the narrow-band filter and write: Assuming a constant comoving density of galaxies inside the narrow-band filter, the number of galaxies is proportional to the comoving volume $V$ \citep[for the explicit expression of the comoving
volume, see][ and references therein]{1999astro.ph..5116H}. |
The final equation is: The mean value of A. can be directly. computed from last equation. | The final equation is: The mean value of $\lambda_z$ can be directly computed from last equation. |
The mean values computed with this method depend slishtlv on the cosmology parametrization selected (except on the value of the IIubble constant) and on the emission line chosen. | The mean values computed with this method depend slightly on the cosmology parametrization selected (except on the value of the Hubble constant) and on the emission line chosen. |
The narrow-bancl filter lavout used determines the approach used to approximate the continuum near (he emission line. | The narrow-band filter layout used determines the approach used to approximate the continuum near the emission line. |
In a general case. we have » fillers and » equations like Eq. 6.. | In a general case, we have $n$ filters and $n$ equations like Eq. \ref{eq:flujo1}, |
one for each filter. | one for each filter. |
This provides us with. at most. 2—1 parameters to paranietrize (he continuum ancl one parameter (o characterize the line. fi). | This provides us with, at most, $n - 1$ parameters to parametrize the continuum and one parameter to characterize the line, $\fluxl$ ). |
A power-law continuum (f.x A") can be fitted with at least three bands. | A power-law continuum $\fluxc \propto \lambda^\beta$ ) can be fitted with at least three bands. |
It is even possible (o assume a given. power-law index (for example. the ¢=—1 lor Lyva emitters in 7?)). | It is even possible to assume a given power-law index (for example, the $\beta=-1$ for $\LA$ emitters in \citet{2006ApJ...645L...9A}) ). |
Another possible functional form for the continuum is a polynomial of degree n— 2. | Another possible functional form for the continuum is a polynomial of degree $n - 2$ . |
Most survevs use one πατοραπ filler on (he region of interest aud one or more fillers (e.g.2?2?).. | Most surveys use one narrow-band filter on the region of interest and one or more broad-band filters \citep[e.g.][]{2001A&A...379..798P,2004AJ....127..563H,2006ApJ...638..596A}. |
Other survevs use more several contiguous narrow-band filters ancl a mecdium-band filter as à veto filler (2) or more complicated lavouts (?).. | Other surveys use more several contiguous narrow-band filters and a medium-band filter as a veto filter \citep{2006A&A...455...61W} or more complicated layouts \citep{2006astro.ph..9497H}. |
In the following sections we study (wo different cases of the polynomial functional form. | In the following sections we study two different cases of the polynomial functional form. |
With two broad-band filters. a linear dependency of the continuum flux wilh wavelength can be assumed. | With two broad-band filters, a linear dependency of the continuum flux with wavelength can be assumed. |
In the case of having one broad-band filter. the continuum has a constant value. | In the case of having one broad-band filter, the continuum has a constant value. |
We develop here a general solution for the case of three fillers covering the line. | We develop here a general solution for the case of three filters covering the line. |
Solutions for othercases with more filters can be easily obtainecl. | Solutions for othercases with more filters can be easily obtained. |
and 2 were determined using the techniques described by ον, 7 and 7.. | and \ref{wd_times} were determined using the techniques described by \citet*{wood85}, , \citet{wood86b} and \citet{wood89a}. |
X median filter was used to smooth the data. the derivative of which was then calculated numerically. | A median filter was used to smooth the data, the derivative of which was then calculated numerically. |
X box-car filter was applied to this derivative. and simple searches were made to locate the minimum anc maximum. values of the derivative corresponding to the midpoints of ingress oO; and egress ©. ( | A box-car filter was applied to this derivative, and simple searches were made to locate the minimum and maximum values of the derivative corresponding to the midpoints of ingress $\phi_{i}$ and egress $\phi_{e}$. ( |
In fact this method locates the stecpest part of the ingress and egress. but we would expect these to correspond to the midpoints unless the light distribution is asvmmetrical.) | In fact this method locates the steepest part of the ingress and egress, but we would expect these to correspond to the midpoints unless the light distribution is asymmetrical.) |
Loa bright spot eclipse is also present. care must be taken to ensure that at this stage the ingress ancl egress of the white cdwarl are not confused with those of the bright spot. | If a bright spot eclipse is also present, care must be taken to ensure that at this stage the ingress and egress of the white dwarf are not confused with those of the bright spot. |
Phe eclipse contact phases corresponding to the start and end of the ingress 6j. ó» and the start and. end of the egress ó5. ὧν were determined by locating the points where the derivative differs significantly from a spline fit to the more slowly varving component. | The eclipse contact phases corresponding to the start and end of the ingress $\phi_{1}$ , $\phi_{2}$ and the start and end of the egress $\phi_{3}$, $\phi_{4}$ were determined by locating the points where the derivative differs significantly from a spline fit to the more slowly varying component. |
Once the white dwarf eclipse contact phases have been found. the white dwarf lishteurve can be reconstructed ane subtracted from the overall lighteurve as illustrated in Figure 2.. | Once the white dwarf eclipse contact phases have been found, the white dwarf lightcurve can be reconstructed and subtracted from the overall lightcurve as illustrated in Figure \ref{deconvolution}. |
The out-of-eclipse white dwarf lluxes thus found are given in. Table 2.. | The out-of-eclipse white dwarf fluxes thus found are given in Table \ref{wd_times}. |
The white dwarl Εικ can be used to determine its temperature and distance: see section 6.. | The white dwarf flux can be used to determine its temperature and distance; see section \ref{sys_parameters}. . |
Once this has been done the bright spot eclipse contact. phases (eiven in Table 3)) can be determined by a similar metho (?) and its lighteurve removed from that of the disc eclipse. | Once this has been done the bright spot eclipse contact phases (given in Table \ref{bs_times}) ) can be determined by a similar method \citep{wood89a} and its lightcurve removed from that of the disc eclipse. |
Η successful. this process can be used to determine the brigh spot temperature. | If successful, this process can be used to determine the bright spot temperature. |
Unfortunately we were unsuccessful in our attempts to do this. probably because Liekering hinclerec accurate determination of the bright spot Hux and contac phases. | Unfortunately we were unsuccessful in our attempts to do this, probably because flickering hindered accurate determination of the bright spot flux and contact phases. |
In the cliscussion that follows we use the sullixes “a” anc b to denote white dwarf and.bright spot contact. phases. respectively 0,4; means the mid-ingresspoint of thewhite cdwarl eclipse). | In the discussion that follows we use the suffixes $w$ ' and $b$ ' to denote white dwarf andbright spot contact phases, respectively $\phi_{wi}$ means the mid-ingresspoint of thewhite dwarf eclipse). |
These median offsets are computed in two iterations. | These median offsets are computed in two iterations. |
We first calculate the median fractional offset for each run in bins of 0.0208" in Dec. ie. 120 bins over the width of Stripe 82. or about 1O bins per CCD width. | We first calculate the median fractional offset for each run in bins of $^{\circ}$ in Dec, i.e., 120 bins over the width of Stripe 82, or about 10 bins per CCD width. |
This exercise is designed to correct flatfielding errors for a given run. | This exercise is designed to correct flatfielding errors for a given run. |
Note that these errors would only depend on Dec because the SDSS emplovs a scan camera. and the scan direction for Stripe 82 is in the RA direction. | Note that these errors would only depend on Dec because the SDSS employs a drift-scan camera, and the scan direction for Stripe 82 is in the RA direction. |
After correcting for the declination-dependent offsets. we then re-compute the median fractional flux offsets for each field along a given run (teach SDSS field is 0.15" long in RA). | After correcting for the declination-dependent offsets, we then re-compute the median fractional flux offsets for each field along a given run (each SDSS field is $^{\circ}$ long in RA). |
This additional field-by-field offset corrects for any temporal variations in the photometric zeropoint of a given run. which are due to transparency/extinction changes over the course of a nominally photometric night. | This additional field-by-field offset corrects for any temporal variations in the photometric zeropoint of a given run, which are due to transparency/extinction changes over the course of a nominally photometric night. |
Flux offsets for a certain wave band were only applied to objects assigned to a bin with at least 9 reference catalogue stars in order to guarantee the accuracy of the derived flux offset. | Flux offsets for a certain wave band were only applied to objects assigned to a bin with at least 9 reference catalogue stars in order to guarantee the accuracy of the derived flux offset. |
In practice. this extra restriction only affects the photometry of objects in the wv band. and in other wave bands when the atmospheric transparency is low. | In practice, this extra restriction only affects the photometry of objects in the $u$ band, and in other wave bands when the atmospheric transparency is low. |
We use a photometric calibration ag (see Table 33) to monitor whetjer or not a flux offset has been applied to the photometry of a {σαΠο object at a certain epoch in ü specific wave band. | We use a photometric calibration tag (see Table \ref{tab:lmc}) ) to monitor whether or not a flux offset has been applied to the photometry of a particular object at a certain epoch in a specific wave band. |
The tinal tsObj files used for subsequent analyses therefore have boh these declination-dependent and field-dependent flux offsets removed for most object records. | The final tsObj files used for subsequent analyses therefore have both these declination-dependent and field-dependent flux offsets removed for most object records. |
We tind that for the standard SDSS runs. the fractional flux offset corrections. which we refer to as photometric zeropoints. are about 1] 2'«. which sets the typical scale of these residual errors in he standard SDSS calibration procedures. | We find that for the standard SDSS runs, the fractional flux offset corrections, which we refer to as photometric zeropoints, are about $1-2\%$, which sets the typical scale of these residual errors in the standard SDSS calibration procedures. |
Figure |. shows the Tactional flux offset corrections as a function of SDSS field number (an arbitrary coordinate along RA assigned to image sections from he same run) for a typical photometric run (94) and a typical yhotometric run (5853). | Figure \ref{fig:zp} shows the fractional flux offset corrections as a function of SDSS field number (an arbitrary coordinate along RA assigned to image sections from the same run) for a typical photometric run (94) and a typical non-photometric run (5853). |
The object catalogues (tsObj files? contain quality and type flags for each object record to aid in the selection of "good" measurements and specitic data samples. | The object catalogues (tsObj files) contain quality and type flags for each object record to aid in the selection of “good” measurements and specific data samples. |
In the LMCC. we only accept object records classified as. galaxies/non-PSF-like objects ΟΡ] file ag = 3) or stars/PSF-like objects (tsObj file tag = 6). and the object must have no child objects (tsObj tile tay = 0: Stoughtonetal. (2002))). | In the LMCC, we only accept object records classified as galaxies/non-PSF-like objects (tsObj file tag $=$ 3) or stars/PSF-like objects (tsObj file tag $=$ 6), and the object must have no child objects (tsObj file tag $=$ 0; \citet{sto2002}) ). |
We then require hat an object record satisties all of a set of constraints in at least one wave band. | We then require that an object record satisfies all of a set of constraints in at least one wave band. |
The tirst of these constraints is that a photometric zeropoint. calculated using the method described in Section 2.2. has been applied to the object record. and that the object record has an PSF magnitude (tsObj file tag 5) brighter han 21.5 for the bands mw. g. rand 7. or brighter than 20.5 for the 2 band. | The first of these constraints is that a photometric zeropoint, calculated using the method described in Section 2.2, has been applied to the object record, and that the object record has an PSF magnitude (tsObj file tag ) brighter than 21.5 for the bands $u$, $g$, $r$ and $i$, or brighter than 20.5 for the $z$ band. |
These limits were chosen to ensure that any photometric measurement in the LMCC has a signal-to-noise ratio of at least 5 in at least one wave band. | These limits were chosen to ensure that any photometric measurement in the LMCC has a signal-to-noise ratio of at least 5 in at least one wave band. |
In Table 2.. we list the remaining set of constraints to be satistied in at least one wave band in order for an object record to be included in trw LMCC. | In Table \ref{tab:req}, we list the remaining set of constraints to be satisfied in at least one wave band in order for an object record to be included in the LMCC. |
We apply one final constraint on the quality of an. object record in order to avoid the incusion of cosmic ray events in our catalogue. | We apply one final constraint on the quality of an object record in order to avoid the inclusion of cosmic ray events in our catalogue. |
If an object record s:itisties all of the above constraints in one wave band only. then it is accepted only if the tsObj file tag for that wave band does not contain the hexadecimal bit OxLO00000. Cag name BE.CE the presence of which indicates that the object is possibly a cosmic ray. | If an object record satisfies all of the above constraints in one wave band only, then it is accepted only if the tsObj file tag for that wave band does not contain the hexadecimal bit 0x1000000 (flag name ), the presence of which indicates that the object is possibly a cosmic ray. |
In order to construc the light-motion curves. we processed each run in turn. starting with the 2005 runs which were closely spaced in time. | In order to construct the light-motion curves, we processed each run in turn, starting with the 2005 runs which were closely spaced in time. |
For each object record in the current run satisfying our quality andtype consraints. we used the following algorithm to process the record: | For each object record in the current run satisfying our quality andtype constraints, we used the following algorithm to process the record: |
About half of the sources in NGC 6946 have possible optical counterparts or dilfuse emission. and the remaining half have neither. | About half of the sources in NGC 6946 have possible optical counterparts or diffuse emission, and the remaining half have neither. |
While the sources for which there are no obvious optical counterparts must be deeply enshrouded. and therefore are likely to be extremely voung. it isn't clear whether detection of light in (he optical regime rules out extreme vouth for regions. | While the sources for which there are no obvious optical counterparts must be deeply enshrouded, and therefore are likely to be extremely young, it isn't clear whether detection of light in the optical regime rules out extreme youth for regions. |
We make this tentative statement [or several reasons: (1) in most cases presented in this sample. the identification of an optical counterpart is insecure due (ο (he possible svstematics between (he optical and radio positions. as well as (he relatively large synthesized radio beam-wicth for. M33: (2) in cases where individual stars are resolved. (M33). themunber of stars required. {ο cereale the Lyman continuum flux are not apparent in the optical images. suggesting either a misidentifieation due to pointing uncertaintv. or a number of the individual stars are. in fact. still enshrouded. | We make this tentative statement for several reasons: (1) in most cases presented in this sample, the identification of an optical counterpart is insecure due to the possible systematics between the optical and radio positions, as well as the relatively large synthesized radio beam-width for M33; (2) in cases where individual stars are resolved (M33), the of stars required to create the Lyman continuum flux are not apparent in the optical images, suggesting either a misidentification due to pointing uncertainty, or a number of the individual stars are, in fact, still enshrouded. |
This scenario has actually been observed in the galactie eccomplex W49A where of the stars in the complex appear to have emerged [rom their birth cocoons while the rest of the complex remains enshrouded (Conti&Blum2001) and (3) if the dominant source of opacity is Thompson scattering (ancl not dust). a source wilh a radius of 5 pe could have electron densities as high as n,=10° and still have opacilies as low as τον1. | This scenario has actually been observed in the galactic complex W49A where of the stars in the complex appear to have emerged from their birth cocoons while the rest of the complex remains enshrouded \citep{conti01} ; and (3) if the dominant source of opacity is Thompson scattering (and not dust), a source with a radius of 5 pc could have electron densities as high as $n_e = 10^5$ $^{-3}$ and still have opacities as low as $\tau \sim 1$. |
Therelore. we conclude that the possible identification of optical counterparts does not preclude (he vouth of these objects. | Therefore, we conclude that the possible identification of optical counterparts does not preclude the youth of these objects. |
However. (hie sources lor which there are no optical counterparts are more likely deeply embedded in (heir natal molecular clouds. and therefore extremely voung. | However, the sources for which there are no optical counterparts are more likely deeply embedded in their natal molecular clouds, and therefore extremely young. |
Given the luminosities and radio spectral energy distributions of rreeions. (heir physical parameters such as size and electron density can be estimated. | Given the luminosities and radio spectral energy distributions of regions, their physical parameters such as size and electron density can be estimated. |
Using the analvtical approximation of Mezeer&IIenderson(1l967).. we can estimate (he emission measure. LAL=[n?dl. given an electron temperature. T. the observing frequency. v. and the optical depth at that [requency. | Using the analytical approximation of \citet{mezgerhenderson67}, we can estimate the emission measure, $EM=\int n_e^2 dl$, given an electron temperature, $T$, the observing frequency, $\nu$, and the optical depth at that frequency. |
The positive spectral index for these sources arises from [ree-[ree emission where 7=I. therefore we assume 7=1 as a lower limit. | The positive spectral index for these sources arises from free-free emission where $\tau \gtrsim 1$, therefore we assume $\tau = 1$ as a lower limit. |
T1e electron temperature of Galactic Ls is ivpically 7.=800041000 Ix (e.g...Atflerbachelal.1996).. which we adopt for this estimate. | The electron temperature of Galactic s is typically $T_e = 8000 \pm 1000$ K \citep[e.g.,][]{afflerbach96}, which we adopt for this estimate. |
The resultinge emission measures for each of tje wavelengths used in (his sample range from ~0.05—16xLOS © pe (for 7>| these emission measures will be correspondingly higher). | The resulting emission measures for each of the wavelengths used in this sample range from $\sim 0.05 - 16
\times 10^8$ $^{-6}$ pc (for $\tau > 1$ these emission measures will be correspondingly higher). |
is about tto the west of EGO 44. | is about to the west of EGO 44. |
As the EGO +4 bow shock faces to the east. IRAS 16367-2356 is likely the exciting source of EGO 44. | As the EGO 44 bow shock faces to the east, IRAS 16367-2356 is likely the exciting source of EGO 44. |
Extensive surveys of outflows toward the p Ophiuchi cloud. including LIL objects in the optical and II» 2.12 eemission in (he near-infrared. have been performed. | Extensive surveys of outflows toward the $\rho$ Ophiuchi cloud, including HH objects in the optical and $_2$ 2.12 emission in the near-infrared, have been performed. |
In total. 46 HII objects. including components wi II objects. P 119 II» emission features have be detected in this region (Grossoetal.2001:Wu2002;Gomezοἱ2003:IXhanzadyan2004:Phelps& | In total, 46 HH objects, including components of HH objects, and 119 $_{2}$ emission features have be detected in this region \citep{gro01,wu02,gom03,kha04,phe04}. |
Barsony "m We identified in n IRAC images 31 EGOs that correspond to the known outflows. among which seven EGOs correspond to known IIL objects and 30 EGOs to IH» near-inlrared emission. | We identified in the IRAC images 31 EGOs that correspond to the known outflows, among which seven EGOs correspond to known HH objects and 30 EGOs to $_{2}$ near-infrared emission. |
Figs. | Figs. |
I8. and 19. show the region of EGO OL which has been identified at 2.12. bby Lucasetal.(2003). | \ref{fig18} and \ref{fig19} show the region of EGO 01 which has been identified at 2.12 by \citet{luc08}. |
. EGO OL are three diffuse nebulae in the 2.12 image while il consists of 3 faint knots in the 3.6 aand 4.5 images. | EGO 01 are three diffuse nebulae in the 2.12 image while it consists of 3 faint knots in the 3.6 and 4.5 images. |
Two YSOs. ο 23 and SR 4. are located in the nearby of EGO01. | Two YSOs, GSS 23 and SR 4, are located in the nearby of EGO 01. |
GSS 23 is classified as a weak-Iline T Tauri star (Bouvier&Appenzeller1992). | GSS 23 is classified as a weak-line T Tauri star \citep{bou92}. |
. Baryetal.(2002) and Bitnerοἱal.(2008) detected narrow ll» emission surrounding GSS 23. | \citet{bar02} and \citet{bit08} detected narrow $_{2}$ emission surrounding GSS 23. |
SR 4 is an emission-line star (Struve& and Phelps&Barsony(2004). detected an HIT object. HII 312. to the southeast of SR 4. | SR 4 is an emission-line star \citep{str49} and \citet{phe04} detected an HH object, HH 312, to the southeast of SR 4. |
Lucaselal.(2008) suggested that there are some connection between (he II» emission ancl the outflow driving by VLA1623-243 (see Figs. 24-- | \citet{luc08} suggested that there are some connection between the $_{2}$ emission and the outflow driving by VLA1623-243 (see Figs. \ref{fig24}- - |
26. and discussions on EGOs in that region) on the basis of their locations aud morphology. | \ref{fig26} and discussions on EGOs in that region) on the basis of their locations and morphology. |
Fies. | Figs. |
5-6 have shown the region of EGO 34 which are counterparts of the outflow JxXCG82004]. [05-04 (IXhanzadyanetal.2004). | 5-6 have shown the region of EGO 34 which are counterparts of the outflow [KGS2004] f05-04 \citep{kha04}. |
. IXhanzadyanetal.(2004) and. al.(2005). observed three knots which correspond to EGO 34a-c in their near-inlrared observations. | \citet{kha04} and \citet{smi05} observed three knots which correspond to EGO 34a-c in their near-infrared observations. |
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