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. However. there is no published evidence for a eamma-rav outburst [rom Gl. and the relatively steady. apparent N-rav [lux (Pooley&Rappaport2006) also argues against a transient source. | However, there is no published evidence for a gamma-ray outburst from G1, and the relatively steady apparent X-ray flux \citep{poo06} also argues against a transient source. |
Thus. the only stellar-mass object that might account for the radio ancl X-ray emission. would be a PWN: using the scaling law given by Frail&Scharringhausen(1997).. we find a likely size ol ~10 milliareseconds lor a PWN radio source in Gl. implvine Chat high-sensitivity VLBI observations could distinguish between a PWN and IMDII origin for the radio emission from Gl. | Thus, the only stellar-mass object that might account for the radio and X-ray emission would be a PWN; using the scaling law given by \citet{fra97}, we find a likely size of $\sim 10$ milliarcseconds for a PWN radio source in G1, implying that high-sensitivity VLBI observations could distinguish between a PWN and IMBH origin for the radio emission from G1. |
]|xnowledge of the radio spectrum of GI could provide more clues to the character οἱ the radio emission. although either a PWN or an IMDBII accretion flow might have a flat spectrum. | Knowledge of the radio spectrum of G1 could provide more clues to the character of the radio emission, although either a PWN or an IMBH accretion flow might have a flat spectrum. |
In any case. our 5 Gllz observation simply is not deep enough. | In any case, our 5 GHz observation simply is not deep enough. |
Hf we choose a 2σ upper limit of 30.0 pJv al 4.9 Giz (2σ chosen since we know the position of the 8.4 Gllz source with high accuracy). we derive a spectral index limit of a>—0.12+0.99 (lor 5,ov[". 1o error in spectral index). which has little power to discriminate among models. | If we choose a $2~\sigma$ upper limit of 30.0 $\mu$ Jy at 4.9 GHz $2~\sigma$ chosen since we know the position of the 8.4 GHz source with high accuracy), we derive a spectral index limit of $\alpha > -0.12 \pm 0.99$ (for $S_\nu\propto \nu^{+\alpha}$, $1~\sigma$ error in spectral index), which has little power to discriminate among models. |
The X-ray emission [rom GI may be due to Boncli aceretion on the IMDII. either from ambient cluster gas or rom stellar winds (Poolev&Rappaport2006). | The X-ray emission from G1 may be due to Bondi accretion on the IMBH, either from ambient cluster gas or from stellar winds \citep{poo06}. |
. IIo.Terashima.&Okajima(2003) and Pooley&Rappaport(2006) give approximate relations for the Boudi accretion on an IMDII in a globular cluster: for an ambient density of 0.1 oE7. an ambient speed of 15 km ! for the gas particles relative to the IMDIL and aradiative efficiency ofLOM... the Bondi accretion Iuminosity for the GI IMDII would be ~3xLO eres |. | \citet{ho03} and \citet{poo06} give approximate relations for the Bondi accretion on an IMBH in a globular cluster; for an ambient density of 0.1 $^{-3}$, an ambient speed of 15 km $^{-1}$ for the gas particles relative to the IMBH, and aradiative efficiency of, the Bondi accretion luminosity for the G1 IMBH would be $\sim 3\times 10^{38}$ ergs $^{-1}$. |
The X-ray luminosity of 2xLO" eres | measured by Pooley&Rappaport(2006). thus implies acerelion ab just under of the Bondi rate. | The X-ray luminosity of $2\times 10^{36}$ ergs $^{-1}$ measured by \citet{poo06} thus implies accretion at just under of the Bondi rate. |
Given that Ly/Lyaq710°. a more likely scenario is (hat G1 aceretes at closer to of the Bondi rate but with a radiative elliciency under 1%. | Given that $L_{\rm X}/L_{\rm Edd} \approx 10^{-6}$, a more likely scenario is that G1 accretes at closer to of the Bondi rate but with a radiative efficiency under . |
.. In this context. we note that the radio/X-rav ratio for Gl is logRy> —4.3. which is above the value of —4.5 used (o divide radio-«quiet Irom racio-loud objects | In this context, we note that the radio/X-ray ratio for G1 is $\log R_{\rm X} > -4.3$ , which is above the value of $-4.5$ used to divide radio-quiet from radio-loud objects |
The emission was obtained by integrating over the apparent disk using 64 concentric rings distributed unevenly over the disk and the limb region. | The emission was obtained by integrating over the apparent disk using 64 concentric rings distributed unevenly over the disk and the limb region. |
The variation in the path lengths through the atmosphere were taken into account fully when calculating the radiation transfer for each ring. | The variation in the path lengths through the atmosphere were taken into account fully when calculating the radiation transfer for each ring. |
In our model the total continuum flux emitted by the surface depends purely on the choice of the temperature To, which defines the temperature scale for the temperature profile to be retrieved. | In our model the total continuum flux emitted by the surface depends purely on the choice of the temperature $T_0$, which defines the temperature scale for the temperature profile to be retrieved. |
We have adjusted Το in such a way as to exactly match the total flux predicted by the ‘Mars continuum model’ of ~4230 Jy provided by ? and to match the temperature fall off towards the limb therein by a factor 0.2. | We have adjusted $T_0$ in such a way as to exactly match the total flux predicted by the `Mars continuum model' of $\sim 4230$ Jy provided by \citet{2008Lellouch} and to match the temperature fall off towards the limb therein by a factor 0.2. |
The error of the modelled flux is5%. | The error of the modelled flux is. |
. Absorption coefficients for the CO spectral lines were calculated using the HITRAN 2008 spectral line catalogue, keeping the terrestrial isotopic ratios in it; i.e., = 498.70 and = 89.01. | Absorption coefficients for the CO spectral lines were calculated using the HITRAN 2008 spectral line catalogue, keeping the terrestrial isotopic ratios in it; i.e., = 498.70 and = 89.01. |
? conclude that the deviations of these isotopic ratios on Mars compared to Earth are less than2%. | \citet{2007Icar..192..396K} conclude that the deviations of these isotopic ratios on Mars compared to Earth are less than. |
. However, to account for carbon dioxide instead of air as broadening gas, the broadening parameters provided by the catalogue were multiplied by a factor of 1.4 according to ?.. | However, to account for carbon dioxide instead of air as broadening gas, the broadening parameters provided by the catalogue were multiplied by a factor of 1.4 according to \citet{1982JQSRT..28..409N}. |
To obtain a temperature profile and the mixing ratio of CO we employed Rodger's optimal estimation method (??).. | To obtain a temperature profile and the mixing ratio of CO we employed Rodger's optimal estimation method \citep{1976RvGeo..14..609R,1990JGR....95.5587R}. |
This method uses so-called ‘a priori’ information — in our case the best estimate of the temperature profile and CO mixing ratio profile — which is then updated to allow for a best fit of the spectral line shapes (minimization of y). | This method uses so-called `a priori' information – in our case the best estimate of the temperature profile and CO mixing ratio profile – which is then updated to allow for a best fit of the spectral line shapes (minimization of $\chi^2$ ). |
This a priori profile was taken from our MAOAM general circulation model (?) considering the exact observation date, geometry, and time of the HIFI observations. | This a priori profile was taken from our MAOAM general circulation model \citep{2005JGRE..11011008H} considering the exact observation date, geometry, and time of the HIFI observations. |
The surface pressure averaged over the visible disk was 6.7 hPa. | The surface pressure averaged over the visible disk was 6.7 hPa. |
Compared to temperature profiles derived from the EMCD, we find temperature differences less than 3 K between 3-60 km. | Compared to temperature profiles derived from the EMCD, we find temperature differences less than 3 K between 3–60 km. |
EMCD provides about 5 K higher temperatures below 3 km and 3 to 5 K higher temperatures between 60—80 km. | EMCD provides about 5 K higher temperatures below 3 km and 3 to 5 K higher temperatures between 60–80 km. |
Up to 100 km, both models slowly merge to almost the same temperature. | Up to 100 km, both models slowly merge to almost the same temperature. |
It is worth noting a high degree of coincidence of the averaged temperature profiles over the field of view of the telescope with the two models. | It is worth noting a high degree of coincidence of the averaged temperature profiles over the field of view of the telescope with the two models. |
Even though the one from EMCD represents monthly averaged fields, the one from MAOAM is based on an instantaneous snapshot, and generally the altitude-latitude distributions differ. | Even though the one from EMCD represents monthly averaged fields, the one from MAOAM is based on an instantaneous snapshot, and generally the altitude-latitude distributions differ. |
This increases our degree of confidence in the simulated temperatures. | This increases our degree of confidence in the simulated temperatures. |
A simultaneous fit of the two spectral lines allows retrieving temperature and mixing ratios independently, because of the rather different optical depth of the two lines, where the line is optically thick (τ=6.3 in line centre) and the line being optically thin (r=1.1). | A simultaneous fit of the two spectral lines allows retrieving temperature and mixing ratios independently, because of the rather different optical depth of the two lines, where the line is optically thick $\tau = 6.3$ in line centre) and the line being optically thin $\tau=1.1$ ). |
Figure b] shows the simultaneously fitted spectra of and (and the residuals) after the removal of the baseline ripple. | Figure \ref{fig:model} shows the simultaneously fitted spectra of and (and the residuals) after the removal of the baseline ripple. |
The retrieved CO mixing ratio amounts to 980+100 ppm which is in agreement with the one detected by SPIRE observations during Ls=5? from 6 November 2009 at 20:20 UT (?) under rather similar surface pressure conditions. | The retrieved CO mixing ratio amounts to $980 \pm 100$ ppm which is in agreement with the one detected by SPIRE observations during $L_s
= 5\degr$ from 6 November 2009 at 20:20 UT \citep{2010arXiv1005.4579S}
under rather similar surface pressure conditions. |
The SPIRE value of 900 ppm has been used as an a priori input to the retrieval algorithm. | The SPIRE value of 900 ppm has been used as an a priori input to the retrieval algorithm. |
Figure B| shows the corresponding temperature profile and the averaging kernels. | Figure \ref{fig:retrieval} shows the corresponding temperature profile and the averaging kernels. |
The latter provide information about the sensitivity of the retrieval versus altitude. | The latter provide information about the sensitivity of the retrieval versus altitude. |
Although the contribution of the a priori profile to the retrieved temperature profile is less than below 60 km, it fits quite well to the profiles predicted by the GCMs. | Although the contribution of the a priori profile to the retrieved temperature profile is less than below 60 km, it fits quite well to the profiles predicted by the GCMs. |
However, the differences are least with EMCD near the ground (5 K) and with MAOAM near 60—70 km. | However, the differences are least with EMCD near the ground $\sim 5$ K) and with MAOAM near 60–70 km. |
Nevertheless, the observations provide about 12-15 K lower temperatures near 65 km. | Nevertheless, the observations provide about 12–15 K lower temperatures near 65 km. |
The temperature inversion between 40-60 km predicted by the GCMs should be manifested in an emission feature in the centre of the CO lines, which obviously is not the case. | The temperature inversion between 40–60 km predicted by the GCMs should be manifested in an emission feature in the centre of the CO lines, which obviously is not the case. |
The error of the model continuum flux translates into a roughly shift of the temperature profiles, i.e., all temperatures will be about 10 K higher or lower than the retrieved value. | The error of the model continuum flux translates into a roughly shift of the temperature profiles, i.e., all temperatures will be about 10 K higher or lower than the retrieved value. |
In the first case, the agreement with the model profile is better between 40--80 km (although still without temperature inversion) but worse below 40 km. | In the first case, the agreement with the model profile is better between 40--80 km (although still without temperature inversion) but worse below 40 km. |
In the second case the agreement is worse for all altitudes. | In the second case the agreement is worse for all altitudes. |
This work presents the first simultaneous retrievals of temperature and carbon monoxide in the Martian atmosphere derived from HIFI data. | This work presents the first simultaneous retrievals of temperature and carbon monoxide in the Martian atmosphere derived from HIFI data. |
The temperature profile can be used as an input parameter for determining of concentrations of other gases observed by HIFI during the same period. | The temperature profile can be used as an input parameter for determining of concentrations of other gases observed by HIFI during the same period. |
Future work will include all observed CO transitions in order to better constrain the temperature profile above 60 km and take advantage of the much wider opacity range for retrieving the vertical profile of CO. | Future work will include all observed CO transitions in order to better constrain the temperature profile above 60 km and take advantage of the much wider opacity range for retrieving the vertical profile of CO. |
Lore. + is the adiabatic index of the gas. V4 is the rate of heating per unit volume due to all the heat sources (i.c. Compton heating and photo-heating) and A is the rate of cooling per unit. volume due to all the heat sinks (i.c. AMremsstrahlung cooling anc various atomic processes). mia is the total number clensity of atoms (HI and Lie) and their ions per unit volume. 7 is the temperature of the shell and Ay is Doltzmann's constant. | Here, $\gamma$ is the adiabatic index of the gas, $\Sigma^T$ is the rate of heating per unit volume due to all the heat sources (i.e. Compton heating and photo-heating) and $\Lambda^T$ is the rate of cooling per unit volume due to all the heat sinks (i.e. Bremsstrahlung cooling and various atomic processes), $n_{\rm tot}$ is the total number density of atoms (H and He) and their ions per unit volume, $T$ is the temperature of the shell and $k_{\rm B}$ is Boltzmann's constant. |
ln the above equation the first. term represents adiabatic cooling due to the expansion of he shell. | In the above equation the first term represents adiabatic cooling due to the expansion of the shell. |
The second term accounts for the elfects of changes in the mean atonic mass due to ionization and recombination processes. | The second term accounts for the effects of changes in the mean atomic mass due to ionization and recombination processes. |
The final term accounts for the heating and cooling effects of the various processes that we now discuss Photoionization heats the shell at a rate of where ££ is the energy. of the sampled. photons which is associated with atom/ion number density nj. 0 is the cllective partial photo-ionization cross section (accounting for secondary ionizations) for the ionization stages of LE and Ile. n2,ey is the number density of photons of energy £7. and Lds the ionization potential of 7. | The final term accounts for the heating and cooling effects of the various processes that we now discuss Photoionization heats the shell at a rate of where $E_i$ is the energy of the sampled photons which is associated with atom/ion number density $n_i$, $\sigma^\prime$ is the effective partial photo-ionization cross section (accounting for secondary ionizations) for the ionization stages of H and He, $n_{\gamma(E)}$ is the number density of photons of energy $E$ , and $E_i$ is the ionization potential of $i$. |
In the above. € accounts for heating by secondary electrons and is given by (2): Compton scattering ofCALB hotons from free electrons causes cooling or heating of the gas a arate of (7) where ep is the "Thompson cross section. an ds the radiation constant. Γον is the tem»erature of the CMD at =0. De ds the number density of electrons. per unit volume and m. is trw mass of an electron. | In the above, ${\mathcal E}$ accounts for heating by secondary electrons and is given by \citep{shull85}: Compton scattering of CMB photons from free electrons causes cooling or heating of the gas at a rate of \citep{peebles_recombination_1968}
where $\sigma_{\rm T}$ is the Thompson cross section, $_{\rm R}$ is the radiation constant, $_{\rm CMB}$ is the temperature of the CMB at $z=0$, $n_{\rm e}$ is the number density of electrons per unit volume and ${\rm m}_{\rm e}$ is the mass of an electron. |
For a typical source in our paper. we find that Compton ieating is insignificant. | For a typical source in our paper, we find that Compton heating is insignificant. |
The initial emission rate of ionizing yhotons for a LO’ starburst with a LO? AL. DLE (cletailed in the next section) is 1.3107 photons s+. | The initial emission rate of ionizing photons for a $10^5$ $_\odot$ starburst with a $10^6$ $_\odot$ BH (detailed in the next section) is $\sim 1.3 \times 10^{51}$ photons $^{-1}$. |
The racius o which Compton heating is important (27). for this scenario at.=10 is about 99 pc. | The radius to which Compton heating is important \citep{Ricotti:08} for this scenario at $z=10$ is about 99 pc. |
As we will sec. this is well below he 0.001.1 Alpe scales that are most relevant. for. I-front evolution and 21 em signals in this work (83): thus. Compton reating will not have a significant ellect on our results. | As we will see, this is well below the 0.001–1 Mpc scales that are most relevant for I-front evolution and 21 cm signals in this work 3); thus, Compton heating will not have a significant effect on our results. |
Photon emission due to single electron. recombination cools the shell at a rate where a, is the rate of the recombination processes for its respective atom/ion number densities. n; (2). | Photon emission due to single electron recombination cools the shell at a rate where $\alpha_{\rm r}$ is the rate of the recombination processes for its respective atom/ion number densities, $n_i$ \citep{verner_atomic_1996}. |
Photon emission due to cielectric recombination cools the shell at a rate where o4 is the rate of the recombination process for Ler (22?)).. | Photon emission due to dielectric recombination cools the shell at a rate where $\alpha_{\rm d}$ is the rate of the recombination process for $^{2+}$ \citep{aldrovandi_radiative_1973,shull_ionization_1982,arnaud_updated_1985}. |
Collisional ionization leads to à cooling rate of where o; is the collisional ionization rate coellicient for the respective atom/ion of number density 2; and £7 is the ionization potential of the respective atom/ion. LL. He and Collisional excitation followed by radiative decay. cools the shell at a rate: where oou and ouuagur are the rates of collisional excitations involving Ll and /— respectively. (?).. | Collisional ionization leads to a cooling rate of where $\alpha_{\rm i}$ is the collisional ionization rate coefficient for the respective atom/ion of number density $n_i$ and $E_i$ is the ionization potential of the respective atom/ion, H, He and $^+$ Collisional excitation followed by radiative decay cools the shell at a rate: where $\alpha_{\rm coll H}$ and $\alpha_{\rm coll {He^+}}$ are the rates of collisional excitations involving H and $^+$ respectively \citep{scholz_collisional_1991}. |
Finally. Bremsstrahlung emission cools the shell at a rate Llere. eo is the permittivity of free space and > is the energv-averaged Gaunt factor (?).. | Finally, Bremsstrahlung emission cools the shell at a rate Here, $\epsilon_0$ is the permittivity of free space and $\gamma$ is the energy-averaged Gaunt factor \citep{sutherland_accurate_1998}. |
‘These coupled dillerential equations are solved. numerically using a standard Itunge-Ixutta method. | These coupled differential equations are solved numerically using a standard Runge-Kutta method. |
As noted. earlier. we focus on carly galaxies of tvpical mass ~10° 107 AL. in total mass and of approximate size a few kpe at most. | As noted earlier, we focus on early galaxies of typical mass $\sim 10^8$ $10^{10}$ $_\odot$ in total mass and of approximate size a few kpc at most. |
We therefore perform most of our calculations al >=10. with one caleulation at z=20 for comparison. | We therefore perform most of our calculations at $z = 10$, with one calculation at $z = 20$ for comparison. |
To calculate the feedback [rom a typical QSO/starlorming galaxy at these epochs. we compute the DII mass function at z=10 using data that is publicly available from the Millennium Simulation (?).. | To calculate the feedback from a typical QSO/starforming galaxy at these epochs, we compute the BH mass function at $z = 10$ using data that is publicly available from the Millennium Simulation \citep{springel05}. |
In Figure. 1.. we show the computed. DII mass function at >=10. where we see that a typical quasar is powered. by BLIs in the mass range 107. 10 M. which we use as a baseline for most of the cases considered in this paper. | In Figure \ref{fig:smbh}, we show the computed BH mass function at $z=10$, where we see that a typical quasar is powered by BHs in the mass range $\sim$ $^5$ $^6$ $_\odot$, which we use as a baseline for most of the cases considered in this paper. |
The turnoverin Figure 1. may be partially due to the finite resolution of the simulation itself: in reality. we expect that the mass function should continue to slowly rise to somewhat smaller masses. | The turnoverin Figure \ref{fig:smbh} may be partially due to the finite resolution of the simulation itself; in reality, we expect that the mass function should continue to slowly rise to somewhat smaller masses. |
In our models. the X-rays [rom the stellar populations are minimal. so we consider. cases | In our models, the X-rays from the stellar populations are minimal, so we consider cases |
to the standard predictions. | to the standard predictions. |
Given the observational error bars, one can conclude from Fig.19 and 20 that the models account nicely for the observational constraints on C and N. In figure 21 we plot the [Na/Fe] ratio for the open cluster sample listed in Table 4 as a function of cluster turnoff mass. | Given the observational error bars, one can conclude from \ref{fig:C_N_smiljanic} and \ref{fig:c1213vsCN_smiljanic} that the models account nicely for the observational constraints on C and N. In figure \ref{fig:Na_Fe_smiljanic} we plot the [Na/Fe] ratio for the open cluster sample listed in Table 4 as a function of cluster turnoff mass. |
Note that the observational data were reported to the solar Na value from we assume in the initial composition of ourmodels?. | Note that the observational data were reported to the solar Na value from we assume in the initial composition of our. |
. Both the predictions and observations show a positive correlation between [Na/Fe] values and stellar mass. | Both the predictions and observations show a positive correlation between [Na/Fe] values and stellar mass. |
Rotation-induced mixing leads to an increase of the amount of Na processed to the surface and allows an explanation for the observed dispersion. | Rotation-induced mixing leads to an increase of the amount of Na processed to the surface and allows an explanation for the observed dispersion. |
There is however an offset of about 0.1 dex between the data and the predictions as was already noticed by who compared their observations with standard model predictions by?. | There is however an offset of about 0.1 dex between the data and the predictions as was already noticed by who compared their observations with standard model predictions by. |
. As a matter of fact, very different observational Na abundances for giant stars have been reported in the literature. | As a matter of fact, very different observational Na abundances for giant stars have been reported in the literature. |
As can be seen in Fig. | As can be seen in Fig. |
21 some studies present [Na/Fe] values as high as 40.6 dex(?),, some only a mild overabundance of 40.2 dex 21)mainBodyCitationEnd9042] Hamdani00, and other solar values(?). | \ref{fig:Na_Fe_smiljanic} some studies present [Na/Fe] values as high as +0.6 dex, some only a mild overabundance of +0.2 dex ] , and other solar values. |
. We refer to for a discussion on the possible causes of these discrepancies. | We refer to for a discussion on the possible causes of these discrepancies. |
The present predictions for the stars more massive than ~ 2 Mo are actually in better agreement with the mild overabundance of [Na/Fe] measured by ?. | The present predictions for the stars more massive than $\sim$ 2 $_{\odot}$ are actually in better agreement with the mild overabundance of [Na/Fe] measured by . |
. Finally we show in Fig. | Finally we show in Fig. |
22 the !60/"7O vs !60/!50 for the G and K giants by and?. | \ref{fig:oxygenHarris} the $^{16}$ $^{17}$ O vs $^{16}$ $^{18}$ O for the G and K giants by and. |
. Included in the figure are our predictions. | Included in the figure are our predictions. |
As discussed in 3, thermohaline mixing affects only slightly the !6O/!5O ratio, and leaves 160/170 unaffected; on the other hand, rotation-induced mixing lowers the Ι6Ο/17Ο ratio, and helps account for the lowest '60/'80 values. | As discussed in 3, thermohaline mixing affects only slightly the $^{16}$ $^{18}$ O ratio, and leaves $^{16}$ $^{17}$ O unaffected; on the other hand, rotation-induced mixing lowers the $^{16}$ $^{17}$ O ratio, and helps account for the lowest $^{16}$ $^{18}$ O values. |
Given the large observational uncertainties, the predictions are reasonably consistent with the O isotopic ratios measured in RGB stars. | Given the large observational uncertainties, the predictions are reasonably consistent with the O isotopic ratios measured in RGB stars. |
In the present paper we have investigated the effects of the thermohaline instability induced by *He-burning that sets in above the RGB bump and of rotation-induced mixing on the evolution and chemical properties of low- and mass stars (1 to 4M.) at solar metallicity. | In the present paper we have investigated the effects of the thermohaline instability induced by $^3$ He-burning that sets in above the RGB bump and of rotation-induced mixing on the evolution and chemical properties of low- and intermediate-mass stars (1 to 4 $_{\odot}$ ) at solar metallicity. |
All the stellar models were computed up to the end of the second dredge-up on the early-AGB, and some of them up to the end of the TP-AGB phase. | All the stellar models were computed up to the end of the second dredge-up on the early-AGB, and some of them up to the end of the TP-AGB phase. |
Predictions are compared to data for lithium, !*C/!3C, [N/C], [Na/Fe], !50/"O, and !60/150 in giant stars with well- masses and evolutionary status on the RGB, clump, early-AGB, and planetary nebulae phases. | Predictions are compared to data for lithium, $^{12}$ $^{13}$ C, [N/C], [Na/Fe], $^{16}$ $^{17}$ O, and $^{16}$ $^{18}$ O in giant stars with well-defined masses and evolutionary status on the RGB, clump, early-AGB, and planetary nebulae phases. |
We find that the theoretical and observational behavioursfor these species are in very good agreement over the whole scrutinized mass range. | We find that the theoretical and observational behavioursfor these species are in very good agreement over the whole scrutinized mass range. |
Thermohaline mixing is confirmed to be the main physical process governing the surface abundances | Thermohaline mixing is confirmed to be the main physical process governing the surface abundances |
'e identiied by a centering algorithüu (APPIND) that searcles for three local maxima iu the central liies of 1ie chip. | are identified by a centering algorithm ) that searches for three local maxima in the central lines of the chip. |
The peaks are assuued to be separate by more than 30 pixels aud have au j»proximate width of 80 pixels. | The peaks are assumed to be separated by more than 30 pixels and have an approximate width of 80 pixels. |
Aperture sizes are 'eevaltated by setting the borders at of ille peak intensity of each order. | Aperture sizes are reevaluated by setting the borders at of the peak intensity of each order. |
A traclue algorithm moves in regular five pixel seps ong the dispersion axls. assessing changes in peaς location for each order. leadiug to a description o‘the aperture position. | A tracing algorithm moves in regular five pixel steps along the dispersion axis, assessing changes in peak location for each order, leading to a two-dimensional description of the aperture position. |
Tle aperture tracing fiction. a second o‘ler Legeudre »olvitonal. is fitted to predefined sample regious of the chip that are less allectec by scattered light. | The aperture tracing function, a second order Legendre polynomial, is fitted to predefined sample regions of the chip that are less affected by scattered light. |
Errors in the two-dimensional apertire border definitions are uswilly below tlree pixels. | Errors in the two-dimensional aperture border definitions are usually below three pixels. |
1 order Legendre polsnotdial is it to the speetrtum. which is then normaizecl. | A 30th order Legendre polynomial is fit to the spectrum, which is then normalized. |
Such a Ισ order polynomial is jusifiel by the complex patteru produced by the flat field amp as it passes through tle spectrometer. as shown by figure 3.. | Such a high order polynomial is justified by the complex pattern produced by the flat field lamp as it passes through the spectrometer, as shown by figure \ref{fig:flatj}. |
Artilicial oscillatious at the aperttwes’ limits are ignored after exlractioi. | Artificial oscillations at the apertures' limits are ignored after extraction. |
Typical RMS of the fit is below 5000 ADU. which may seem high but actually amourus to roughly of the average siglla. | Typical RMS of the fit is below 5000 ADU, which may seem high but actually amounts to roughly of the average signal. |
The final flat-field image bas all its pixel counts set to 1. except those on the regious occupied N the spectrum. which are replaced by the ratio between the origial coum aud the fitted polynonial. | The final flat-field image has all its pixel counts set to 1, except those on the regions occupied by the spectrum, which are replaced by the ratio between the original count and the fitted polynomial. |
Inthe NIR spectral regiou the atmosphere plays au important role. | In the NIR spectral region the atmosphere plays an important role. |
Besides a siguilicant telluric absorption. several atinospleric emission lines are entangled with the spectrum o‘the astrouomical source (see figure 1. [or a suuple spectrum of the sky. where the J. Η aucl Ix baids are identified). | Besides a significant telluric absorption, several atmospheric emission lines are entangled with the spectrum of the astronomical source (see figure \ref{fig:sky} for a sample spectrum of the sky, where the J, H and K bands are identified). |
The process of removit© telluric emission lines is conunonuly kuown as sky subtraction. or sky chopping. aud the augular size of the targe dictates whether additional off-sot106 exposures are required. | The process of removing telluric emission lines is commonly known as sky subtraction, or sky chopping, and the angular size of the target dictates whether additional off-source exposures are required. |
Iu the case of poiib sources. whicl OCCUDV only a small fraction of the slit. oue cau ake exposures will the source in two cliΠοιοί positions along the slit. and later subtract susenient images. | In the case of point sources, which occupy only a small fraction of the slit, one can take exposures with the source in two different positions along the slit, and later subtract subsequent images. |
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