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During eeress (he EW rises to ~70A. as the thaw increases. exceeding even its out-of-eclipse value.
During egress the EW rises to $\sim$ 70, as the flux increases, exceeding even its out-of-eclipse value.
In all cases. a double-peaked profile is observed.
In all cases, a double-peaked profile is observed.
The emission line profile during mid-eclipse extends to + 200 kin/s with sienilicantly less [Iux in the extended wines. as compared to the profile during egress. which extends to c 300 km/s. Figure 3 shows the corresponding I) line profiles.
The emission line profile during mid-eclipse extends to $\pm$ 200 km/s with significantly less flux in the extended wings, as compared to the profile during egress, which extends to $\pm$ 300 km/s. Figure 3 shows the corresponding $\beta$ line profiles.
Again. dramatic EW and line profile changes are evident.
Again, dramatic EW and line profile changes are evident.
Although the o1t-ol-eclipse I> line profile is heavily affected by the underlying stellar absorption spectrum. it is similar to Ho.
Although the out-of-eclipse $\beta$ line profile is heavily affected by the underlying stellar absorption spectrum, it is similar to $\alpha$.
In the egress spectrum. however. {he emission occurs primarily on the blue side with a wing extending to -300 km/s. The asvmmnmetirve in this line is much more striking (han Wa.
In the egress spectrum, however, the emission occurs primarily on the blue side with a wing extending to -300 km/s. The asymmetry in this line is much more striking than $\alpha$.
During mid-eclipse. alihough little 132 [lux is present. the shape of the line profile appears similar to eeress.
During mid-eclipse, although little $\beta$ flux is present, the shape of the line profile appears similar to egress.
The weak features visible in this profile at - 25 km/s are probably due to an improper background subtraction of the IL? emission from NGC! 2264.
The weak features visible in this profile at - 25 km/s are probably due to an improper background subtraction of the $\beta$ emission from NGC 2264.
Figure 4 shows the emission-line profiles for the [Ol] 6300 line.
Figure 4 shows the emission-line profiles for the [OI] 6300 line.
The [OI] line. which is weak. but clearly discernable at maximum light. becomes prominent at mid-eclipse and during egress.
The [OI] line, which is weak, but clearly discernable at maximum light, becomes prominent at mid-eclipse and during egress.
The line (lux seems (o be about the same out of eclipse and during eegress. alihough slightlv higher at mid-eclipse.
The line flux seems to be about the same out of eclipse and during egress, although slightly higher at mid-eclipse.
We caution against anv extreme interpretation of (his measurement.
We caution against any extreme interpretation of this measurement.
The profile during mid-eclipse was disturbed by an improper background subiraction. and we [eel that our errors are most likely underestimated.
The profile during mid-eclipse was disturbed by an improper background subtraction, and we feel that our errors are most likely underestimated.
However. this behavior would indicate Chat none of the [OI] emitting zone sulfers variable occultation a (he phases of our observations.
However, this behavior would indicate that none of the [OI] emitting zone suffers variable occultation at the phases of our observations.
This is fully expected given the spatial extent of the (bipolar) forbidden line emitting regions in CI'T5s.
This is fully expected given the spatial extent of the (bipolar) forbidden line emitting regions in CTTSs.
The peak of the high velocity component of the OI] emission in typical CTTSs originates at about 30 AU from the star (Hirthetal.1997 )..
The peak of the high velocity component of the [OI] emission in typical CTTSs originates at about 30 AU from the star \citep{Hirth97}. .
the literature.
the literature.
To scale these SFRs and masses to the ? initial mass function they should be multiplied by a factor of 0.66.
To scale these SFRs and masses to the \citet{Chabrier2003} initial mass function they should be multiplied by a factor of 0.66.
Most of the stellar mass within kkpe of the radio galaxy is concentrated in the object A. which presumably is the radio host galaxy.
Most of the stellar mass within kpc of the radio galaxy is concentrated in the object A, which presumably is the radio host galaxy.
If we only consider the eemitters. more than of the stellar mass lies within the clump labelled A. The instantaneous star formation is more spread out. although object A still contributes of the total SFR.
If we only consider the emitters, more than of the stellar mass lies within the clump labelled A. The instantaneous star formation is more spread out, although object A still contributes of the total SFR.
This central concentration of mass and widespread star formation is reminiscent of 11138-262 (2)..
This central concentration of mass and widespread star formation is reminiscent of 1138-262 \citep{Hatch2009}.
We must be cautious about the mass and SFR estimates for object A as this object may contain an active galactic nucleus (AGN) that can contribute to both the aand continuum luminosity.
We must be cautious about the mass and SFR estimates for object A as this object may contain an active galactic nucleus (AGN) that can contribute to both the and continuum luminosity.
However. object A is significantly extended in both the Ks and VB images. and there is no evidence of point-like emission from a central unobscured AGN.
However, object A is significantly extended in both the $Ks$ and $NB$ images, and there is no evidence of point-like emission from a central unobscured AGN.
So it is unlikely that there is a large contribution of light from the AGN.
So it is unlikely that there is a large contribution of light from the AGN.
However. a partially obscured AGN could still contaminate the Ks and ΠαΠιιλος, and our stellar mass and SFR estimate of this object could be too high.
However, a partially obscured AGN could still contaminate the $Ks$ and fluxes, and our stellar mass and SFR estimate of this object could be too high.
The eemitters near the radio galaxy are aligned along the radio jets ata position angle of .
The emitters near the radio galaxy are aligned along the radio jets at a position angle of $^{\circ}$.
A Monte Carlo simulation of the distribution of the eemitters results in only a chance that the eemitters will be randomly orientated in a cone within of the radio jets.
A Monte Carlo simulation of the distribution of the emitters results in only a chance that the emitters will be randomly orientated in a cone within $^{\circ}$ of the radio jets.
Therefore their alignment with the radio jets imply a causal connection.
Therefore their alignment with the radio jets imply a causal connection.
Radio triggered star formation has been observed in several nearby and distant radio galaxies (e.g??)..
Radio triggered star formation has been observed in several nearby and distant radio galaxies \cite[e.g][]{vanBreugel1985,Bicknell2000}.
The alignment of the eemitters with the radio jets suggest this system may be a good candidate for such jet-triggered star formation.
The alignment of the emitters with the radio jets suggest this system may be a good candidate for such jet-triggered star formation.
However the emitters may also be material that is ionized by the central AGN (see ? for a review).
However the emitters may also be material that is ionized by the central AGN (see \citealt{MileydeBreuck2008} for a review).
In order to determine whether these are dwarf galaxies or simply clumps of ionized gas we searched for stellar continuum in other wavebands.
In order to determine whether these are dwarf galaxies or simply clumps of ionized gas we searched for stellar continuum in other wavebands.
Objects A. B and D have relatively strong Keon continuum fluxes. likely emitted from an underlying stellar population.
Objects A, B and D have relatively strong $K_{\rm cont}$ continuum fluxes, likely emitted from an underlying stellar population.
The stellar masses of these objects range from 10? πο [0M...
The stellar masses of these objects range from $10^{9}$ to $10^{11}$.
Object A and its nearest neighbour D lie at approximately the same redshift (2). and are possibly interacting or Merging as they are joint by a single emission line halo.
Object A and its nearest neighbour D lie at approximately the same redshift \citep{Iwamuro2003}, and are possibly interacting or merging as they are joint by a single emission line halo.
Object C is detected in both Fo06W (25.6 mag) and FI60W. mmag) in the Hubble Space Telescope (HS7) images of ??.. but object E is not detected in either image.
Object C is detected in both F606W (25.6 mag) and F160W mag) in the Hubble Space Telescope ) images of \citet{Pentericci1999, Pentericci2001}, but object E is not detected in either image.
Nebular emission (both continuum and emission lines) contributes to the flux in each of theseEST bands so it is difficult to unambiguously determine whether any of the emission is stellar continuum.
Nebular emission (both continuum and emission lines) contributes to the flux in each of these bands so it is difficult to unambiguously determine whether any of the emission is stellar continuum.
We conclude that objects A. B and D are galaxies. whose star formation may be triggered or enhanced by the radio jets.
We conclude that objects A, B and D are galaxies, whose star formation may be triggered or enhanced by the radio jets.
Object D is very bright inΠα.. and has a similar SFR to clump A. however it contains a relatively small amount of mass.
Object D is very bright in, and has a similar SFR to clump A, however it contains a relatively small amount of mass.
If stars were continuously forming in object D at its current rate. its stellar mass could be formed in only MMyrs. which is a similar timescale to the lifetime of the radio emission.
If stars were continuously forming in object D at its current rate, its stellar mass could be formed in only Myrs, which is a similar timescale to the lifetime of the radio emission.
The nature of objects C and E is unclear. they may be dwarf galaxies with masses less than 109?M... or simply pockets of dense gas ionized by the AGN or undergoing their first burst of star formation.
The nature of objects C and E is unclear, they may be dwarf galaxies with masses less than $10^{9.5}$, or simply pockets of dense gas ionized by the AGN or undergoing their first burst of star formation.
However. the location of these objects suggest that the radio jets are responsible for the enhanced eemission that makes them visible.
However, the location of these objects suggest that the radio jets are responsible for the enhanced emission that makes them visible.
The massive radio galaxy may be surrounded by several dwarf galaxies. or pockets of gas. that cannot be detected. but may become visible when their line emission increases as they pass through the radio jets.
The massive radio galaxy may be surrounded by several dwarf galaxies, or pockets of gas, that cannot be detected, but may become visible when their line emission increases as they pass through the radio jets.
In this section we compare the density of eemitters in the radio galaxy fields to the control. fields.
In this section we compare the density of emitters in the radio galaxy fields to the control fields.
1138-262 lies in a denseproto-cluster.. with many spectroscopically confirmed eemitting members (7?)..
1138-262 lies in a dense, with many spectroscopically confirmed emitting members \citep{Kurk2004a,Kurk2004b}.
If aalso lies within a proto-clusters.— it should be surrounded by a large overdensity of eemitters. of comparable density to the ΤΙ135-262cluster.
If also lies within a s, it should be surrounded by a large overdensity of emitters, of comparable density to the 1138-262.
The surface density of gealaxies per NB magnitude is the most basic observed quantity which can determine whether a region is overdense or not.
The surface density of galaxies per $NB$ magnitude is the most basic observed quantity which can determine whether a region is overdense or not.
The uncertainties and significance of this measurement can be wel determined from a counts-in-cells analysis of the control field.
The uncertainties and significance of this measurement can be well determined from a counts-in-cells analysis of the control field.
Five of the line-emitters in the ftield are influenced by the radio galaxy —(see Section 3.1.29). whils 3 galaxies lie within the Ένα halo of 11138-262. so may also be influenced by the radio source in this field.
Five of the line-emitters in the field are influenced by the radio galaxy (see Section \ref{distribution}) ), whilst 3 galaxies lie within the $\alpha$ halo of 1138-262, so may also be influenced by the radio source in this field.
We are interested in the large-scale structure surrounding these radio galaxies. beyond the immediate influence of the radio Jets. so these sources were removed from the catalogues and not included in the remainder of this study.
We are interested in the large-scale structure surrounding these radio galaxies, beyond the immediate influence of the radio jets, so these sources were removed from the catalogues and not included in the remainder of this study.
The surface densities of ccandidates surrounding the two radio galaxies and 1138-262 are shown in refoverdensity.. where they are compared to the control field.
The surface densities of candidates surrounding the two radio galaxies and 1138-262 are shown in \\ref{overdensity}, where they are compared to the control field.
The density of oobjects in the GOODS-S field was adjusted to account for the difference in narrow-band filter bandwidths.
The density of objects in the GOODS-S field was adjusted to account for the difference in narrow-band filter bandwidths.
The quoted uncertainties are 16 standard deviations resulting from a counts-in-cells analysis measured from 6.9 or aaremin? cells in the control fields.
The quoted uncertainties are $\sigma$ standard deviations resulting from a counts-in-cells analysis measured from 6.9 or $^2$ cells in the control fields.
These are the sizes of the aud [1138-262 fields. respectively.
These are the sizes of the and 1138-262 fields respectively.
The Ισ standard
The $\sigma$ standard
performed at 1.6 GHz for two hours with the EVN only at 512 Mbps data rate.
performed at 1.6 GHz for two hours with the EVN only at 512 Mbps data rate.
These observations served as a pilot project designed to assess the feasibility of the experiment; we refer to them as epoch 0.
These observations served as a pilot project designed to assess the feasibility of the experiment; we refer to them as epoch $0$.
The other three experiments (epochs 1,2,3) were carried out with a global array at 22 GHz for maximum resolution, although Australia-Europe baselines could not be formed at the third epoch because of the availability constraints of the Australian telescopes.
The other three experiments (epochs $1,2,3$ ) were carried out with a global array at 22 GHz for maximum resolution, although Australia-Europe baselines could not be formed at the third epoch because of the availability constraints of the Australian telescopes.
Each observation lasted for about 11 hours and was preceded by a 2-3 hour preparation and clock-search run for the Australian-Asian array, and later included an hour of clock-search and re-referencing for the whole array including the EVN.
Each observation lasted for about 11 hours and was preceded by a 2-3 hour preparation and clock-search run for the Australian-Asian array, and later included an hour of clock-search and re-referencing for the whole array including the EVN.
The fringe-finders used were 0420-014, OJ287, and 1055+018.
The fringe-finders used were $-$ 014, OJ287, and 1055+018.
Additional nearby calibrators 0736+017 and 0805—077 were regularly observed for continuous fringe monitoring; these later also served as amplitude calibration check sources.
Additional nearby calibrators 0736+017 and $-$ 077 were regularly observed for continuous fringe monitoring; these later also served as amplitude calibration check sources.
The total flux densities of these compact sources were obtained from single-dish Effelsberg (2nd epoch) and synthesis-array ATCA (3rd epoch) measurements.
The total flux densities of these compact sources were obtained from single-dish Effelsberg (2nd epoch) and synthesis-array ATCA (3rd epoch) measurements.
The total bit rate per telescope was 512 Mbps, divided into four 16 MHz sub-bands in both polarizations.
The total bit rate per telescope was 512 Mbps, divided into four 16 MHz sub-bands in both polarizations.
The data were streamed in real-time from the stations to the EVN data processor (Schilizzietal.2001) at JIVE in all experiments.
The data were streamed in real-time from the stations to the EVN data processor \citep{Schilizzi2001} at JIVE in all experiments.
Plots of the real-time fringes as well as the final coverage of the (u, v)-plane are shown in Giroletti (2010c)..
Plots of the real-time fringes as well as the final coverage of the $(u,v)$ -plane are shown in \citet{Giroletti2010c}.
The longest baseline in our array was the on June 10, reaching kkm.
The longest baseline in our array was the Yebes--Mopra on June 10, reaching km.
The data were processed using the NRAO Astronomical Image Processing System (AIPS).
The data were processed using the NRAO Astronomical Image Processing System (AIPS).
Amplitude calibration was done with the measured gain curves and systemtemperatures.
Amplitude calibration was done with the measured gain curves and systemtemperatures.
For the LBA stations, nominal SEFDs were used.
For the LBA stations, nominal SEFDs were used.
The Τεν. data were corrected for atmospheric opacity assuming standard weather conditions for a subset of stations that had opacity-free calibration data.
The $T_{\rm sys}$ data were corrected for atmospheric opacity assuming standard weather conditions for a subset of stations that had opacity-free calibration data.
This in general worked well except for the lowest elevations below 10°.
This in general worked well except for the lowest elevations below $^\circ$.
All the sources were fringe-fit separately, no phase-referencing being applied.
All the sources were fringe-fit separately, no phase-referencing being applied.
All of the stations produced fringes, but Kashima had setup problems in the first two experiments and Hobart had a setup problem in the first part of the first experiment.
All of the stations produced fringes, but Kashima had setup problems in the first two experiments and Hobart had a setup problem in the first part of the first experiment.
After fringe-fitting, the station amplitudes were further adjusted using the calibrator sources.
After fringe-fitting, the station amplitudes were further adjusted using the calibrator sources.
Phase and amplitude self-calibration were performed in Difmap to produce the final images.
Phase and amplitude self-calibration were performed in Difmap to produce the final images.
The phase response of the right- and left-handed polarized signals with parallactic angle was corrected before fringe-fitting for all stations, including the new 40m Yebes radio telescope.
The phase response of the right- and left-handed polarized signals with parallactic angle was corrected before fringe-fitting for all stations, including the new 40m Yebes radio telescope.
For this Nasmyth offset mount antenna, we used the routines of the version of AIPS, following the procedures described in Dodson (2009); no RR/LL phase
For this Nasmyth offset mount antenna, we used the routines of the version of AIPS, following the procedures described in \citet{Dodson2009}; ; no RR/LL phase
be zero outside this range.
be zero outside this range.
Our criterion is based on examining the effects of cutting the number counts at a given level on a selection of galaxy evolution models from the literature.
Our criterion is based on examining the effects of cutting the number counts at a given level on a selection of galaxy evolution models from the literature.
We compare the predicted ffor each literature model truncated below a specified flux level with the wwithout truncation. and tind that our data are not sensitive to a cut-off of less than 0.1 mJy at 250um.
We compare the predicted for each literature model truncated below a specified flux level with the without truncation, and find that our data are not sensitive to a cut-off of less than 0.1 mJy at 250.
A similar analysis shows that truncating the tit above | Jy is also undetectable. with similar values for the other passbands.
A similar analysis shows that truncating the fit above 1 Jy is also undetectable, with similar values for the other passbands.
From simulations. we find that we can obtain good constraints if the second knot lies approximately at the lo instrumental noise.
From simulations, we find that we can obtain good constraints if the second knot lies approximately at the $1 \sigma$ instrumental noise.
Because the number counts below our flux limit are unlikely to be well described by a single knot all the way down to 0.1 mJy. the fit value for this point should be treated with care: simulations indicate that this is not a problem for the 2? mJy knot.
Because the number counts below our flux limit are unlikely to be well described by a single knot all the way down to 0.1 mJy, the fit value for this point should be treated with care; simulations indicate that this is not a problem for the 2 mJy knot.
In order to avoid over-tuning our fits to represent literature models. we adopt approximately logarithmically spaced knots between these extremes.
In order to avoid over-tuning our fits to represent literature models, we adopt approximately logarithmically spaced knots between these extremes.
The choice of the number of knots is somewhat arbitrary.
The choice of the number of knots is somewhat arbitrary.
Neighboring knots are very strongly correlated. and as the number increases the correlations increase.
Neighboring knots are very strongly correlated, and as the number increases the correlations increase.
We have tried to chose the number of knots to be as large as possible while keeping the correlations reasonably small.
We have tried to chose the number of knots to be as large as possible while keeping the correlations reasonably small.
We fit all three fields simultanously. but each band independentut
We fit all three fields simultanously, but each band independently.
The uncertainty in the instrumental noise is modeled as a single multiplicative factor having a Gaussian prior with σ=SG,
The uncertainty in the instrumental noise is modeled as a single multiplicative factor having a Gaussian prior with $\sigma = 5\%$.
Note that we are making the assumption that the timestream instrumental noise is the same for all three fields as found in Nguyen (2010).
Note that we are making the assumption that the timestream instrumental noise is the same for all three fields as found in \citet{Nguyen:2010}.
. In addition to the SPIRE data. we also explore the ettects of including the FIRAS CFIRB prior (Fixsenetal.1998). by integrating SdN/dS for our model down to the lowest knot and adding a term to the likelihood that compares that value with the FIRAS measurement and its error.
In addition to the SPIRE data, we also explore the effects of including the FIRAS CFIRB prior \citep{Fixsen:1998} by integrating $S dN/dS$ for our model down to the lowest knot and adding a term to the likelihood that compares that value with the FIRAS measurement and its error.
This assumes that the CFIRB is entirely due to discrete sources. and that flux densities outside the range of our model contribute only negligibly.
This assumes that the CFIRB is entirely due to discrete sources, and that flux densities outside the range of our model contribute only negligibly.
We integrate the Fixsenetal.(1998). spectrum through the SPIRE passbands and adopt the relative errors given in Marsdenetal.(2009).
We integrate the \citet{Fixsen:1998} spectrum through the SPIRE passbands and adopt the relative errors given in \citet{mar09}.
The uncertainty in the relative calibration between FIRAS and SPIRE significantly affects the utility of this prior.
The uncertainty in the relative calibration between FIRAS and SPIRE significantly affects the utility of this prior.
The best fit multiply broken power-law fit is compared with the GOODS-N data in figure 3.. and the parameters are given in tables 2. and 3. and for the spline interpolation fits in tables + and 5..
The best fit multiply broken power-law fit is compared with the GOODS-N data in figure \ref{fig:fitplot}, , and the parameters are given in tables \ref{tbl:brokpownoprior} and \ref{tbl:brokpowprior}, and for the spline interpolation fits in tables \ref{tbl:splinenoprior} and \ref{tbl:splineprior}.
The correlations between adjacent knots are large andnegative"... with typical correlation coefficients of —0.5 to —-0.8.
The correlations between adjacent knots are large and, with typical correlation coefficients of $-0.5$ to $-0.8$.
The twomodels are compared with each other in figure 4.
The twomodels are compared with each other in figure \ref{fig:comparemodel}.
The two interpolating models (spline and multiply-broken power-law) produce very similar results.
The two interpolating models (spline and multiply-broken power-law) produce very similar results.
As discussed previously. these are model fits. not independent number counts. and since the parameterizations differ. directly comparing the values at the knot positions is not entirely correct.
As discussed previously, these are model fits, not independent number counts, and since the parameterizations differ, directly comparing the values at the knot positions is not entirely correct.
Nonetheless. the agreement is clear.
Nonetheless, the agreement is clear.