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We found that the optical excess of the DD-baud flux (Zane et al. | We found that the optical excess of the B-band flux (Zane et al. |
2008) with respect to the best- Naav spectra extrapolation is now 21.1411.1 (30 confidence level). | 2008) with respect to the best-fit X-ray spectrum extrapolation is now $24.4 \pm 11.1$ $3 \sigma$ confidence level). |
This value is still incompatible with rotation-powered non-thermal emission fromL. unless its optical eiuissiou efficieucy is ~1000 times larecr iui that of rotatiou-powered pulsars (Zharikov ot al. | This value is still incompatible with rotation-powered non-thermal emission from, unless its optical emission efficiency is $\sim 1000$ times larger than that of rotation-powered pulsars (Zharikov et al. |
20063. | 2006). |
Aloreover. our R-baud upper Im would 1nost ikely nuplv a PL specral iudex a<Q0. while rotation-oowered pulsars have usually 0xaxd (AGenani ct al. | Moreover, our R-band upper limit would most likely imply a PL spectral index $\alpha < 0$, while rotation-powered pulsars have usually $0\leq \alpha \leq 1$ (Mignani et al. |
20075: 010...0), | 2007b; 2010a,b). |
Ou the other hae. explaining the yptical excess in terms of pure thermal emission roni the jeutron star surface would require both a small distance aud rather striugeut Bits on the inclination of he LOS aud the maeuetic axis with respect to the neutron star spin axis. | On the other hand, explaining the optical excess in terms of pure thermal emission from the neutron star surface would require both a small distance and rather stringent limits on the inclination of the LOS and the magnetic axis with respect to the neutron star spin axis. |
For instance. a radius rpx15 kan nuplies a neutron star distance dx200 pe. which implies Το20 eV aud ry1.3 kun | For instance, a radius $r_O \leq 15$ km implies a neutron star distance $d\leq 200$ pc, which implies $T_O \leq 20$ eV and $r_X \leq 1.3$ km. |
Such a small Nav cluitting area ds compatible with the 0.01 N-rav pulsed fraction if either the star spin axis is ucarly aligned with the imagnetic axis or with the LOS. or it is slightly misaligned with respect to both the maeuctic axis aud the LOS bv 5. 107. | Such a small X-ray emitting area is compatible with the $\sim 0.04$ X-ray pulsed fraction if either the star spin axis is nearly aligned with the magnetic axis or with the LOS, or it is slightly misaligned with respect to both the magnetic axis and the LOS by $5-10^{\circ}$ . |
Observations iu the near-UV would obviously be crucial to constrain the slope of the ooptical spectrmm and to better constrain Te) aud ro. | Observations in the near-UV would obviously be crucial to constrain the slope of the optical spectrum and to better constrain $T_O$ and $r_O$. |
With the fluxes falling well above all possible R-J spectra compatible with the flux measurements. ouly the (HEST)) can provide the required neu-UV sensitivity. | With the fluxes falling well above all possible R-J spectra compatible with the flux measurements, only the ) can provide the required near-UV sensitivity. |
Future optical observations. both from the erouud aud from space. will also be iumportaut to measure the source proper motion. to independently confirm: the optical identification and tointer dudirect constraints on the distance. | Future optical observations, both from the ground and from space, will also be important to measure the source proper motion, to independently confirm the optical identification and toinfer indirect constraints on the distance. |
That estimate is also consistent with the conclusion by ?,, who estimated that for 6<z«9, which translates into Ny,S2.5 if the contribution of sources beyond z=9 is unimportant. | That estimate is also consistent with the conclusion by \citet{m03}, who estimated that for $6<z<9$, which translates into $N_{\gamma/b}\la2.5$ if the contribution of sources beyond $z=9$ is unimportant. |
Thus, a requirement appears to be a sensible criterion for the reionization of the universe by z=6. | Thus, a requirement appears to be a sensible criterion for the reionization of the universe by $z=6$. |
The same condition has also been obtained by ? from extrapolating the production rate of ionizing photons required to fit the observed evolution of the mean opacity of the Lyman-alpha forest to z=6. | The same condition has also been obtained by \citet{bh07} from extrapolating the production rate of ionizing photons required to fit the observed evolution of the mean opacity of the Lyman-alpha forest to $z=6$. |
Thus, the WMAPI cosmology is marginally sufficient to satisfy the condition (4)), while the WMAP3 universe is well short of the needed amount of radiation at z6 by at least a factor of 2 (and, perhaps, as ionizingmuch as a factor of 10 if the value for the intrinsic break rin, is close to 6 than to 3, and the transition to low escape fraction occurs at 5x10!!Mc rather than at 5x10/9 M). | Thus, the WMAP1 cosmology is marginally sufficient to satisfy the condition \ref{eq:crit}) ), while the WMAP3 universe is well short of the needed amount of ionizing radiation at $z\approx6$ by at least a factor of 2 (and, perhaps, as much as a factor of 10 if the value for the intrinsic break $r_{\rm int}$ is close to 6 than to 3, and the transition to low escape fraction occurs at $5\times10^{11}\Msun$ rather than at $5\times10^{10}\Msun$ ). |
This, somewhat unexpected, result crucially depends on the main conclusion of ? that the escape fraction is very small for low mass galaxies. | This, somewhat unexpected, result crucially depends on the main conclusion of \citet{gkc07} that the escape fraction is very small for low mass galaxies. |
That conclusion is consistent with the observational measurements of the escape fraction by ? and our knowledge of dwarf galaxies in the local universe, which are known to have large gas fractions and eextend that exceeds the extend of the stellar disk. | That conclusion is consistent with the observational measurements of the escape fraction by \citet{flc03} and our knowledge of dwarf galaxies in the local universe, which are known to have large gas fractions and extend that exceeds the extend of the stellar disk. |
On the other hand, as Figure 4 shows, if the escape fraction is independent of the galaxy mass (smin= 1), the WMAP3 cosmology comfortably falls into the reionization requirement with N./,(z=6)2:1.5. | On the other hand, as Figure \ref{figMIN} shows, if the escape fraction is independent of the galaxy mass $smin=1$ ), the WMAP3 cosmology comfortably falls into the reionization requirement with $N_{\gamma/b}(z=6)\approx1.5$. |
It is, therefore,imperative to have the measurements of the escape fraction extended to even fainter galaxies and result of ? verified with higher resolution simulations and different numerical methods. | It is, therefore,imperative to have the measurements of the escape fraction extended to even fainter galaxies and result of \citet{gkc07} verified with higher resolution simulations and different numerical methods. |
Were it found that the escape fractions of dwarf galaxies are, indeed, negligibly small, then new, more exotic sources of ionizing radiation (Pop III stars, X-ray binaries, a new, previously unknown population of faint quasars, etc)) would need to be invoked to explain the (relatively) early reionization of the universe at z2” 6. | Were it found that the escape fractions of dwarf galaxies are, indeed, negligibly small, then new, more exotic sources of ionizing radiation (Pop III stars, X-ray binaries, a new, previously unknown population of faint quasars, ) would need to be invoked to explain the (relatively) early reionization of the universe at $z\approx 6$ . |
as green lines in Figure 1, includes two relatively strong SO, lines within the o-H5O spectrum, and a weaker line within the p-H2O spectrum. | as green lines in Figure 1, includes two relatively strong $_2$ lines within the $_2$ O spectrum, and a weaker line within the $_2$ O spectrum. |
These lines are included, as an additional background, in the calculation of the H5O optical depth. | These lines are included, as an additional background, in the calculation of the $_2$ O optical depth. |
Another SO; line is present in the o-H;8O spectrum at -15 kmss~! (Fig. | Another $_2$ line is present in the $_2^{18}$ O spectrum at $-15$ $^{-1}$ (Fig. |
2, upper panel) and has been similarly modelled. | 2, upper panel) and has been similarly modelled. |
The full source model of S.-L. Qin (private comm.) | The full source model of S.-L. Qin (private comm.) |
suggests that no other strong lines are present in our spectra at velocities corresponding to the foreground clouds. | suggests that no other strong lines are present in our spectra at velocities corresponding to the foreground clouds. |
Figure 2 (upper and middle panels) shows the o- and H,°O and H;°0 spectra divided by the background emission, including dust continuum and the SO; lines. | Figure 2 (upper and middle panels) shows the o- and $_2^{16}$ O and $_2^{18}$ O spectra divided by the background emission, including dust continuum and the $_2$ lines. |
The o-H}°O is an equally-weighted average of the band 1a and 1b data. | The $_2^{16}$ O is an equally-weighted average of the band 1a and 1b data. |
The lines are saturated over a wide range of velocities and thus H,°Ounusable for a quantitative analysis. | The $_2^{16}$ O lines are saturated over a wide range of velocities and thus unusable for a quantitative analysis. |
However, we have identified several velocity ranges with moderate saturation levels, marked with thick horizontal lines in Figure 2. | However, we have identified several velocity ranges with moderate saturation levels, marked with thick horizontal lines in Figure 2. |
These can be identified with the expanding molecular ring —50 55” !), a transition between the 4 kpc arm and the Orion arm (—12 to —7 kmss!), the Sagittarius arm (5 to 20 ss!), and the Scutum arm (27 to 35 kmss'!; possibly blended with the Sagittarius B2 envelope; see e.g. Neufeld et al. | These can be identified with the expanding molecular ring $<-50$ $^{-1}$ ), a transition between the 4 kpc arm and the Orion arm $-12$ to $-7$ $^{-1}$ ), the Sagittarius arm (5 to 20 $^{-1}$ ), and the Scutum arm (27 to 35 $^{-1}$; possibly blended with the Sagittarius B2 envelope; see e.g. Neufeld et al. |
2000). | 2000). |
Assuming that the foreground absorption completely covers the continuum and all water molecules are in the ground state (a reasonable assumption for the diffuse foreground clouds given the very high spontaneous emission rate coefficients for the ground-state water lines, 3.458x107? and 1.842x107? s! for the ortho and para lines, respectively), we derive optical depths of the o- and p-water lines (r=—In //J,). | Assuming that the foreground absorption completely covers the continuum and all water molecules are in the ground state (a reasonable assumption for the diffuse foreground clouds given the very high spontaneous emission rate coefficients for the ground-state water lines, $3.458 \times 10^{-3}$ and $1.842
\times 10^{-2}$ $^{-1}$ for the ortho and para lines, respectively), we derive optical depths of the o- and p-water lines $\tau = -\ln I/I_o$ ). |
The resulting optical depth ratio is shown in Figure 2 (lower panel; left intensity scale). | The resulting optical depth ratio is shown in Figure 2 (lower panel; left intensity scale). |
An ortho/para optical depth ratio of 1 corresponds to a column density ratio of 2. | An ortho/para optical depth ratio of 1 corresponds to a column density ratio of 2. |
The resulting ortho/para column density ratio is given by the right hand scale in Figure 2 (lower panel). | The resulting ortho/para column density ratio is given by the right hand scale in Figure 2 (lower panel). |
The uncertainty in the line optical depth in a given velocity channel is given by ór=exp(r)x6//I,. | The uncertainty in the line optical depth in a given velocity channel is given by $\delta \tau = \exp (\tau) \times \delta I /I_o$. |
The errorbars in Figure 2 (lower panel), are computed under the conservative assumption ó7/1I,=0.05 (maximum uncertainty, dominated by the standing waves, based on a comparison of the continuum normalized spectra of the o-H2O line in bands 1a and 1b; the signal-to-noise ratio for the p-H5O line is comparable, given the stronger continuum, and the amplitude of the standing waves also scales with the continuum strength). | The errorbars in Figure 2 (lower panel), are computed under the conservative assumption $\delta I/I_o = 0.05$ (maximum uncertainty, dominated by the standing waves, based on a comparison of the continuum normalized spectra of the $_2$ O line in bands 1a and 1b; the signal-to-noise ratio for the $_2$ O line is comparable, given the stronger continuum, and the amplitude of the standing waves also scales with the continuum strength). |
Relative uncertainties of the o- and p-water optical depths have been added in quadrature. | Relative uncertainties of the o- and p-water optical depths have been added in quadrature. |
Table 1 gives weighted averages of the ortho/para column density ratio in different velocity ranges towards Sagittarius B2(M), along with the corresponding uncertainties. | Table 1 gives weighted averages of the ortho/para column density ratio in different velocity ranges towards Sagittarius B2(M), along with the corresponding uncertainties. |
Since the individual measurements are not independent but dominated by instrumental systematics, we use a more | Since the individual measurements are not independent but dominated by instrumental systematics, we use a more |
The region around the dynamic center of our Galaxy is very active recently in massive star formation (e.g.. Figer et al. | The region around the dynamic center of our Galaxy is very active recently in massive star formation (e.g., Figer et al. |
1999). which should have vielded various high-energy products such as supernova remnants (SNRs) and neutron stars (e.g. Morris Serabyn 1996: Cordes Lazio 1997: Wane. Gotthelf Lang 2002). | 1999), which should have yielded various high-energy products such as supernova remnants (SNRs) and neutron stars (e.g., Morris Serabyn 1996; Cordes Lazio 1997; Wang, Gotthelf Lang 2002). |
Young and fast-rotating neutron | Young and fast-rotating neutron |
We present four models of hot Jupiter circulation here. two thal are drag-[ree ancl two that include a simple treatment of magnetic drag in the atmosphere. | We present four models of hot Jupiter circulation here, two that are drag-free and two that include a simple treatment of magnetic drag in the atmosphere. |
The cdrae-free models differ in (hat one uses (he same hyperdissipation strength from ?.. while the other uses a hvperdissipalion Umescale an order of magnitude shorter in order (o reduce the nunmerical noise (hat is apparent al low pressures in the original model. | The drag-free models differ in that one uses the same hyperdissipation strength from \citet{rau10}, while the other uses a hyperdissipation timescale an order of magnitude shorter in order to reduce the numerical noise that is apparent at low pressures in the original model. |
The two magnetic drag models differ in the implementation of drag at low pressures. as discussed below. | The two magnetic drag models differ in the implementation of drag at low pressures, as discussed below. |
The (hree-climensional models analvzed in (his paper are all basic extensions of models presented previously in ? and ?.. and details can be found in those papers. | The three-dimensional models analyzed in this paper are all basic extensions of models presented previously in \citet{rau10} and \citet{per10a}, and details can be found in those papers. |
Drieflv. (μον were all caleulated. using the same dynamical core and Newtonian relaxation scheme for radiative forcing as presented in ?.. | Briefly, they were all calculated using the same dynamical core and Newtonian relaxation scheme for radiative forcing as presented in \citet{men09}. |
Each model was run for 2000 planet days at a resolution of T31L45 (corresponding to a horizontal resolution of ~4 and 45 vertical levels in log pressure). | Each model was run for 2000 planet days at a resolution of T31L45 (corresponding to a horizontal resolution of $\sim 4^\circ$ and 45 vertical levels in log pressure). |
Two modifications have been made (o the set-up of these models in order (o better [acilitate calculation of transmission spectra: a different specifie gas constant was used and the upper boundary has been extended up to LO ματ. | Two modifications have been made to the set-up of these models in order to better facilitate calculation of transmission spectra: a different specific gas constant was used and the upper boundary has been extended up to 10$\mu$ bar. |
The models [from ?. and 2. used a specific gas constant of R=4593 J ! !. which corresponds to a mean molecular mass of 1.51 g/mol. and a value lor Hey of 0.321 (cy is the specilic heat capacity at constant pressure). | The models from \citet{rau10} and \citet{per10a} used a specific gas constant of $R=4593$ J $^{-1}$ $^{-1}$, which corresponds to a mean molecular mass of 1.81 g/mol, and a value for $R/c_p$ of 0.321 $c_p$ is the specific heat capacity at constant pressure). |
For better consistency wilh the transmission spectroscopy modeling. which is hiehly sensitive to atmospheric scale height and therefore mean molecular weight (e.g.?). we chose to use 2=3523 J ! kg! and Rie,=0.286 in (he circulation models presented here. | For better consistency with the transmission spectroscopy modeling, which is highly sensitive to atmospheric scale height and therefore mean molecular weight \citep[e.g.][]{mil09} , we chose to use $R=3523$ J $^{-1}$ $^{-1}$ and $R/c_p =0.286$ in the circulation models presented here. |
This now corresponds to a mean molecular weight of 2.36 g/mol. a value more appropriate [or solar composition. while the value for {όν now malches that for a purely diatomic gas. | This now corresponds to a mean molecular weight of 2.36 g/mol, a value more appropriate for solar composition, while the value for $R/c_p$ now matches that for a purely diatomic gas. |
The other significant change over the previously published models is that we have extended the upper boundary of the atmosphere to 10 αχ. necessary because transmission spectra probe pressures well above the 1 mbar top boundary of the original models. | The other significant change over the previously published models is that we have extended the upper boundary of the atmosphere to 10 $\mu$ bar, necessary because transmission spectra probe pressures well above the 1 mbar top boundary of the original models. |
This extension lo lower pressures neant (hat we had to extrapolate our radiative forcing and crag prescriptions. | This extension to lower pressures meant that we had to extrapolate our radiative forcing and drag prescriptions. |
The Newtonian relaxation scheme used for the radiative forcing requires the choice of equilibrium temperatures and radiative times (lordetailssee2).. | The Newtonian relaxation scheme used for the radiative forcing requires the choice of equilibrium temperatures and radiative times \citep[for details see][]{rau10}. |
In order to extend (hese up to 10 jrbar we chose to continue the 1000 Ix. dav-night temperature difference. which is constant down to LOO mbar. | In order to extend these up to 10 $\mu$ bar we chose to continue the 1000 K day-night temperature difference, which is constant down to 100 mbar. |
The night side temperature was then set in the same manner as in our previous work. so that at each level the integrated 7! matched the prolile from Figure 1 of ?.. | The night side temperature was then set in the same manner as in our previous work, so that at each level the integrated $T^4$ matched the profile from Figure 1 of \citet{iro05}. |
The radiative times were taken from Figure 4 of ?.. whieh includes these lower | The radiative times were taken from Figure 4 of \citet{iro05}, , which includes these lower |
A major characteristic time scale in X-ray binaries is their orbital period. which. in many cases. is seen in their light curves as periodic modulations. see. e.g.. Wenetal.(2006). for a list of orbital periodicities from the All-Sky Monitor (ASM) aboard (RXTE: Bradt.Rothschild&Swank1993:Levineetal. 1996)). | A major characteristic time scale in X-ray binaries is their orbital period, which, in many cases, is seen in their light curves as periodic modulations, see, e.g., \citet{wen06} for a list of orbital periodicities from the All-Sky Monitor (ASM) aboard ; \citealt*{brs93,levine96}) ). |
The observed flux modulations appear due to a number of ditferent physical effects. see Poutanen.Zdziarski&Ibragimov(2008) thereafter PZIOS) for a list of the possibilities. | The observed flux modulations appear due to a number of different physical effects, see \citet*{pzi08} (hereafter PZI08) for a list of the possibilities. |
In addition. a number of X-ray binaries show modulation at quasi-periods much longer than their orbital periods. | In addition, a number of X-ray binaries show modulation at quasi-periods much longer than their orbital periods. |
Those period modulations. listed. e.g.. in Wenetal.(2006) and Ogilvie&Dubus (2001).. hereafter ΟΡΟΙ. are called superorbital (hereafter SO). | Those long-period modulations, listed, e.g., in \citet{wen06} and \citet{ogdu01}, hereafter OD01, are called superorbital (hereafter SO). |
In most of the known cases. the SO periodicity (or quasi-periodicity) appears to be caused by precession of an accretion disc and/or jet. which either results in variable obscuration of emitted X-rays as in Her X-1 (Katz1973:Roberts1974).. or in changes the viewing angle of the presumed anisotropic emitter. as in SS 433 (Katz1980) or Cyg X-I (e.g. Lachowiezetal.2006.. hereafter 106. Ibragimov.Zdziarski&Poutanen2007.. hereafter IZPO7). or both. | In most of the known cases, the SO periodicity (or quasi-periodicity) appears to be caused by precession of an accretion disc and/or jet, which either results in variable obscuration of emitted X-rays as in Her X-1 \citep{k73,roberts74}, or in changes the viewing angle of the presumed anisotropic emitter, as in SS 433 \citep{k80} or Cyg X-1 (e.g., \citealt*{l06}, hereafter L06, \citealt*{i07}, hereafter IZP07), or both. |
In some relatively rare cases. the SO periodicity is caused by modulation of the accretion rate. e.g.. in the low-mass X-ray binaries 4U [820—303 (Zdziarski.Wen&Gierlifiski20072). and 4U 1636-536 (Shihetal.2005:Farrell.Barret&Skinner 2009).. | In some relatively rare cases, the SO periodicity is caused by modulation of the accretion rate, e.g., in the low-mass X-ray binaries 4U 1820–303 \citep*{z07a} and 4U 1636-536 \citep*{shih05,fbs09}. . |
In particular. Cyg X-1. a well-known high-mass X-ray binary containing a persistent accreting black-hole source (see. e.g. Ziólkowski2005 for the properties of the binary). has shown such (quasi) periodicity in its X-ray. optical and radio emission with an average SO period of =150 d. found in a large number of investigations (BrocksoppP,etal.1999a:Pooley.Fender&Brock-&Demircan2001:Benllochetal.2001. 2004:: L06: IZPO7). | In particular, Cyg X-1, a well-known high-mass X-ray binary containing a persistent accreting black-hole source (see, e.g., \citealt{zi05} for the properties of the binary), has shown such (quasi) periodicity in its X-ray, optical and radio emission with an average SO period of $\psup\simeq 150$ d, found in a large number of investigations \citealt*{brocksopp99a,poo99,kitamoto00,kar01,od01,benlloch01,benlloch04}; L06; IZP07). |
The modulation is detectable in the hard X-ray spectral state (see. e.g. Zdziarski&Gierlinski 2004). in which Cyg X-1 spends most of the time. | The modulation is detectable in the hard X-ray spectral state (see, e.g., \citealt{zg04}) ), in which Cyg X-1 spends most of the time. |
No SO modulation has been detected in the soft states. but this seems to be due to their typicalshort duration and strong overall variability. | No SO modulation has been detected in the soft states, but this seems to be due to their typicalshort duration and strong overall variability. |
The underlying physical process appears to still | The underlying physical process appears to still |
Understanding the physics of convective core overshoot is important in interpretiig the color-2magnitude diagrams (CMD8) aud the Iuminosityv functions of open star clusters. ancl eenerallv of intermediate-age stellar populations. | Understanding the physics of convective core overshoot is important in interpreting the color-magnitude diagrams (CMD's) and the luminosity functions of open star clusters, and generally of intermediate-age stellar populations. |
This is due to the facts that increasiig the mixed core mass mocdilies the evolutionary tracks of stars (allecting the isochrone swpe). and lengthens the evolutionary lifetimes of stars (affecting age determinations and huninoslly [functions near (he main sequence turnoff). | This is due to the facts that increasing the mixed core mass modifies the evolutionary tracks of stars (affecting the isochrone shape), and lengthens the evolutionary lifetimes of stars (affecting age determinations and luminosity functions near the main sequence turnoff). |
Thus. in addition to the classical probem of | Thus, in addition to the classical problem of |
INSs whose spectra are softer 2011). | INSs whose spectra are softer . |
. It is possible that the optical/UV spectral iudex is related to the magnetic field. either directly through the magnetosphere or indirectly through shifting spectral lines 2011): in that case the simnmiluitv of aand wwould be natural. | It is possible that the optical/UV spectral index is related to the magnetic field, either directly through the magnetosphere or indirectly through shifting spectral lines ; in that case the similarity of and would be natural. |
Overall. our mieasuromeut firmly places aas one of the INSs despite its short period. aud moves us sigmificautly closer to having a complete sampled of measured. spin-downs for that population. | Overall, our measurement firmly places as one of the INSs despite its short period, and moves us significantly closer to having a complete sampled of measured spin-downs for that population. |
There are still a umber of open questions to be answered via X-rav aud multiwavelength observations. | There are still a number of open questions to be answered via X-ray and multi-wavelength observations. |
Primary among these is understanding the surface cussion through consistent modeling of the spectra and lghtcurves. aud ideally with phase-resolved spectroscopy. | Primary among these is understanding the surface emission through consistent modeling of the spectra and lightcurves, and ideally with phase-resolved spectroscopy. |
Observations at optical/UV wavelengths of the pulsed. emission could make significant imurprovenienuts in our uuderstaudiug. bv tviug the emitting areas at different wavelcneths together and establishing the deeree of surface imhomoecueicty. | Observations at optical/UV wavelengths of the pulsed emission could make significant improvements in our understanding, by tying the emitting areas at different wavelengths together and establishing the degree of surface inhomogeneiety. |
Finally. further kinematic ages would help ereatlv in constraining the coupled evolution of maeguctic field aud temperature. | Finally, further kinematic ages would help greatly in constraining the coupled evolution of magnetic field and temperature. |
Based ὃν observations obtained with NAILNewton. an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. | Based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and NASA. |
DLE was partially supported by NASA through erant NNNOSANSOG. Apart from the NMMSAS data reduction pipelines provided bvNewton. this rescarch has made use of software provided bv the Chandra X-ray Center (CNC) in the application packages CTAO aud Sherpa. | DLK was partially supported by NASA through grant NNX08AX39G. Apart from the XMMSAS data reduction pipelines provided by, this research has made use of software provided by the Chandra X-ray Center (CXC) in the application packages CIAO and Sherpa. |
below. | below. |
The linmnua continu emission ds detected at Soscup —2.040.22 indy (Figure. 6. right). iu good agrecnuent with the extrapolation from the SCUBA lneasurelent {and consistent with the estimate of Tacconi et 2006. of Sosscui 523:0.5 nud. their Table 1). | The mm continuum emission is detected at $_{\rm
234\,GHz}$ $\pm$ 0.22 mJy (Figure 6, right), in good agreement with the extrapolation from the SCUBA measurement (and consistent with the estimate of Tacconi et 2006, of $_{\rm
233\,GHz}$ $\pm$ mJy, their Table 1). |
We have observed both the aand eenission in 6 additional sources. four quasars and two submillimeter galaxies. | We have observed both the and emission in 6 additional sources, four quasars and two submillimeter galaxies. |
All fitted parameters aud line fluxes are sunuuarized in Table 3. | All fitted parameters and line fluxes are summarized in Table 3. |
In the following we briefly discuss the individual sources. | In the following we briefly discuss the individual sources. |
This source is oue of the (apparentlv) brightest sources known to date. with a subimillameter fux of (Lewis ct 11998). | This source is one of the (apparently) brightest sources known to date, with a submillimeter flux of $_{850\mu m}$ $\pm$ mJy (Lewis et 1998). |
The first CO detections=ot by SssujanLumdyDownes et ((1999. J=l. J=9) already tudicated the extreme excitation conditious of the molecular gas (studied in detail in et 22007). | The first CO detections by Downes et (1999, $J$ =4, $J$ =9) already indicated the extreme excitation conditions of the molecular gas (studied in detail in et 2007). |
One of the explanations for its apparent brigltuess aud extreme excitation is that the leusine ecometry is such that the central ppc of this source is magnified by a large factor (μου.LOO et 22007. sec also Egaii et22000)?. | One of the explanations for its apparent brightness and extreme excitation is that the lensing geometry is such that the central pc of this source is magnified by a large factor $\mu$ =60–100 et 2007, see also Egami et. |
. The Ine has been detected by Wage et ((2006). | The line has been detected by Wagg et (2006). |
The source is not detected in eenissiou with an upper liuit of 1JJyvlhaunss! (Figure 7. left: upper lits are given in Table 3). | The source is not detected in emission with an upper limit of $^{-1}$ (Figure 7, left; upper limits are given in Table 3). |
This source is bright in the submillimeter Darvainis Ivison 2002) and CO has 726.7beenEl (SssomΕΤ.tentatively detected in this source by Uainhue ct ((2001. 7223). | This source is bright in the submillimeter $_{850\mu
m}$ $\pm$ mJy, Barvainis Ivison 2002) and CO has been tentatively detected in this source by Hainline et (2004, $J$ =3). |
Followup observations by Wl (.7=3.5.7.8.9) show that the CO enissiou line is very narrow (FWIAI~Li0kaus 1)). | Follow–up observations by W11 $J$ =3,5,7,8,9) show that the CO emission line is very narrow $\sim$ ). |
We detect both the aand eenission lines in this source (see Figure 8 and Table 3). | We detect both the and emission lines in this source (see Figure 8 and Table 3). |
DRII13535O17 is bright in the submillimeter (6.9. Beutord et 11999. we adopt 8559, 7224 ΕΙΝ from their Figure 2) and has first been detected in CO (4-5) bv Caullotean et (1997). | 1335–0417 is bright in the submillimeter (e.g., Benford et 1999, we adopt $_{850\mu m}$ $\pm$ mJy from their Figure 2) and has first been detected in CO $J$ =5) by Guilloteau et (1997). |
[heh aneular resolution CO observations revealed a couples. possibly mereine system (Carll et 22002a. Riechers et 22008a. J=2). | High angular resolution CO observations revealed a complex, possibly merging system (Carilli et 2002a, Riechers et 2008a, $J$ =2). |
The source is not detected in both the aand liue with upper limits of aud ως| (Figure 9. right: see also Table 3). | The source is not detected in both the and line with upper limits of $^{-1}$ and $^{-1}$ (Figure 9, right; see also Table 3). |
The spectrumJJyv of the Ine in Figure 9 also includes the CO(76) line that is close in frequency space (the CO emission is discussed iu WII). | The spectrum of the line in Figure 9 also includes the CO(7–6) line that is close in frequency space (the CO emission is discussed in W11). |
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