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The stellar mass added to the bulge during the burst increases the bulge to total stellar mass ratio. and the cold gas supply is depleted. | The stellar mass added to the bulge during the burst increases the bulge to total stellar mass ratio, and the cold gas supply is depleted. |
Around z~5.2 there is a another burst of star formation. triggered by a major merger. which leads to a large increasein the bulge-to-total ratio of the galaxy from 0.3 to 0.9. | Around $z \sim 5.2$ there is a another burst of star formation, triggered by a major merger, which leads to a large increasein the bulge-to-total ratio of the galaxy from 0.3 to 0.9. |
However. there is less star formation associated with this burst. due to the depletion of the gas reservoir prior to the burst.and the most massive progenitor is only a faint LBG. | However, there is less star formation associated with this burst, due to the depletion of the gas reservoir prior to the burst,and the most massive progenitor is only a faint LBG. |
Another burst happens at 2o 3.5. which makes the most massive progenitor galaxy a faint LBG but with relatively little change in the 2/7 ratio. | Another burst happens at $z \sim 3.5$ , which makes the most massive progenitor galaxy a faint LBG but with relatively little change in the $B/T$ ratio. |
The 5/7 ratio gradually declines as quiescent star formation adds mass to | The $B/T$ ratio gradually declines as quiescent star formation adds mass to |
uaenitude pw1.2% at a position angele of ()zzas’, | magnitude $p\approx1.2\%$ at a position angle of $\theta\approx88\degr$. |
We row demonstrate how such au instruuental polarization nav be corrected for. | We now demonstrate how such an instrumental polarization may be corrected for. |
The fact that polarized standards have bxο Ποσο Oo be consistent with previous studies (CTIA-DC-F7. TD2s3812) alows us to be confident that the results or the uupoarized standards are real o:uid that we nay be seeing a residual iustruineutal polarization. | The fact that polarized standards have been measured to be consistent with previous studies (CHA-DC-F7, HD283812) allows us to be confident that the results for the unpolarized standards are real, and that we may be seeing a residual instrumental polarization. |
The ow nunber of observed uipolarized standards and the oedestal effec make it difficult to be confident iu the exact nature of the iustimental poluization. but we can nevertheless make au ateimpt to correct the data for his. | The low number of observed unpolarized standards and the pedestal effect make it difficult to be confident in the exact nature of the instrumental polarization, but we can nevertheless make an attempt to correct the data for this. |
As the degree aud orientation of polirizafion are derived in Stokes (£.Q.U) space this is also where we can perform an instrumeutal correction. | As the degree and orientation of polarization are derived in Stokes $I,Q,U$ ) space this is also where we can perform an instrumental correction. |
It is esseutial hat the correction is carried out iu the (Q. Uj-lane for each standard and not iu the celestial coordinate svstein. | It is essential that the correction is carried out in the $Q,U$ )-plane for each standard and not in the celestial coordinate system. |
The unpolarized standard data from pre- and post-NCS observations are considered separatelv. | The unpolarized standard data from pre- and post-NCS observations are considered separately. |
While the (Q.U) quadrauts are cousistent between the )olinized standards. the actual values of Q aud C are dependent ou J£. | While the $Q,U$ ) quadrants are consistent between the polarized standards, the actual values of $Q$ and $U$ are dependent on $I$. |
The corrections are therefore averaged (weighted by the errors iu Q aud C) across all observed staudards after beiug normalized by the intensity. rather than being simply carried out iu the (Q.C )-planc | The corrections are therefore averaged (weighted by the errors in $Q$ and $U$ ) across all observed standards after being normalized by the intensity, rather than being simply carried out in the $Q,U$ )-plane. |
«The results from the correction are presented in Table 5.. | The results from the correction are presented in Table \ref{tab:correct}. |
It can be seen that the «Ispersion around the expected aud previously reporte results has been reduced. | It can be seen that the dispersion around the expected and previously reported results has been reduced. |
TheY 47D for the polarized sandards about the expected values is now ~0.01 as compared to ~0.1 withotit the (QU) correction. | The $\chi^2$ for the polarized standards about the expected values is now $\sim0.01$ as compared to $\sim0.1$ without the $Q,U$ ) correction. |
Tt is also possible o null the polarization ii1 uipolarized targets by adjusting the on orbit derived transmission | It is also possible to null the polarization in unpolarized targets by adjusting the on orbit derived transmission co-efficients. |
As Q = U = 0 in unpolarized sources. fj can be derived by considering J,=LY, aud fixing ty=0.9667. | As Q = U = 0 in unpolarized sources, $t_k$ can be derived by considering $I_k=IX_k$ and fixing $t_3 = 0.9667$. |
Solving f4 aud fo for p—0 and 0=07 in IID331891 across all photometric apertures results in tl1ο co-efüicieut profiles preseuted in Figure L.. | Solving $t_1$ and $t_2$ for $p=0\%$ and $\theta=0\degr$ in HD331891 across all photometric apertures results in the co-efficient profiles presented in Figure \ref{fig:coeffs}. |
It can be see iin Fieure | that in the central most aperture the co-ctticj0nts derived from the 2002 epoch (dashed line) are cosistent with the co-cficicuts quoted in Table 2.. | It can be seen in Figure \ref{fig:coeffs} that in the central most aperture the co-efficients derived from the 2002 epoch (dashed line) are consistent with the co-efficients quoted in Table \ref{tab:coeffs}. |
Similar stall scale variations to that of the observed polarimeric curves of erowthn are also seen. | Similar small scale variations to that of the observed polarimetric curves of growth are also seen. |
Taking au average o the derived 2003 epoch co-effiicieuts. outside of the innero wc-xecond. and carrving through the photometric errOrs. GIVES US f,=STILT£0.0005 and to=0.53410.005. | Taking an average of the derived 2003 epoch co-efficients, outside of the inner arc-second, and carrying through the photometric errors, gives us $t_1 = 0.8717\pm0.0005$ and $t_2 = 0.8341\pm0.0005$. |
Applviug these adjusted co-cfiicicuts Gvhich are within of the original co-effücieuts) to the 2003 data for CTIA-DC-F7 produces a typical change of dp50. imid 00zz3° over the valies of p aud 0 derived. from the original efficicuts. | Applying these adjusted co-efficients (which are within of the original co-efficients) to the 2003 data for CHA-DC-F7 produces a typical change of $\delta{p}\approx0.1\%$ and $\delta{\theta}\approx3\degr$ over the values of $p$ and $\theta$ derived from the original co-efficients. |
Adjusting the rausnussion co-cficicuts to null the measured polavization iu a single uupoluidzed standard lay provide a “¢wick fix” for some data. but it is by πο lucas a concrete solution to the underlying problem. | Adjusting the transmission co-efficients to null the measured polarization in a single unpolarized standard may provide a “quick fix” for some data, but it is by no means a concrete solution to the underlying problem. |
After all. it is possible there have bee1 no changes in the frauswission co-cticicuts. | After all, it is possible there have been no changes in the transmission co-efficients. |
Ideally we would be able to characterize these effects by obscrving several polarized. standards at the cardinal angles of je polarizers (this greatv reduces the degrees of freedou iu the coefficieut matrix). | Ideally we would be able to characterize these effects by observing several polarized standards at the cardinal angles of the polarizers (this greatly reduces the degrees of freedom in the coefficient matrix). |
This would also allow us to see if there is an instrmucutal polarization that is dependent on the orientatio raud level of polarization im the object as well as the ro] augle ofZLST. | This would also allow us to see if there is an instrumental polarization that is dependent on the orientation and level of polarization in the object as well as the roll angle of. |
For exiuude. it is possible that the sensitivity of the mirrors o polarized light is producing some reflective. retardaice. | For example, it is possible that the sensitivity of the mirrors to polarized light is producing some reflective retardance. |
Unfortunately. there is. a present. insutiicicut calibraion data available in the areuve with which to fully check this possibility. | Unfortunately, there is, at present, insufficient calibration data available in the archive with which to fully check this possibility. |
There does exist data for two poluized objects. CITÀ-DC-F? aib DID2823812. at amultiple roll angles (sec Table 1)). but this statistically iusignificautf sample inhibits 1s from drawing anv conclusious about the nature of an observationally depeudoeut lustrumental polarization. | There does exist data for two polarized objects, CHA-DC-F7 and HD283812, at multiple roll angles (see Table \ref{tab:sample}) ), but this statistically insignificant sample inhibits us from drawing any conclusions about the nature of an observationally dependent instrumental polarization. |
While xevious polarimetric studies with NICMOS have not needed to be concerned witli suc1 characteristics due to the high levels of observed volarization (Ueta.Aburahmwwa&Meisner 2005). as he capabilities. of NICAOS ire stretchecd bv. more and nire anabitious programs so will tre need to accurately measure low levels of polarizaticπα. | While previous polarimetric studies with NICMOS have not needed to be concerned with such characteristics due to the high levels of observed polarization \citep{umandm05}, as the capabilities of NICMOS are stretched by more and more ambitious programs so will the need to accurately measure low levels of polarization. |
Iu fact. a munber of programs necding these leveIs of accuracy have already Όσοι executed. | In fact, a number of programs needing these levels of accuracy have already been executed. |
For exaure. 110160 (The imclear scattering ecolctry of Sevtert ealaxics} aud LOLLO (Anisotropy and obscuration iu the near πιο]ar reeious of powerful radio galaxies) will voth require Ligh accuracies as AGN eenerally display pml5. | For example, 10160 (The nuclear scattering geometry of Seyfert galaxies) and 10410 (Anisotropy and obscuration in the near nuclear regions of powerful radio galaxies) will both require high accuracies as AGN generally display $p\approx1-5\%$. |
Iu additio Leas tje Jaunes Webb Space Telescope will uo be diving with poαποαν optics. NICMOS wi Lremain the oilv iustraucut capable of performing such high precision Πασάς polarimetry on faint objects. | In addition, as the James Webb Space Telescope will not be flying with polarimetry optics, NICMOS will remain the only instrument capable of performing such high precision imaging polarimetry on faint objects. |
The receut polarization calibrajon for the Advanced Camera for Survey".u (Diettactal.2001:Diretta& prexfs ano eoxaluXe for the NICMOS | The recent polarization calibration for the Advanced Camera for Surveys \citep{bir04,bandkp04,kpandb04,kpandb05} presents an example for the NICMOS |
Ες aud a noise component Fa. | $F_S$ and a noise component $F_N$. |
Ideally. we require that the detector [Iux change. OF. is entirely due to the stellar RV change dey. | Ideally, we require that the detector flux change, $\delta F$, is entirely due to the stellar RV change $\delta v_S$. |
However. 9£ is also partly induced by telluric liue shift dea: resulting [rom raudom atmospheric motious. | However, $\delta F$ is also partly induced by telluric line shift $\delta v_N$ resulting from random atmospheric motions. |
Therefore. both dey aud ovy contribute to δΕ | Therefore, both $\delta v_S$ and $\delta v_N$ contribute to $\delta F$. |
We have two sets oL RV measurements. ovg+o(0.btpyys.s) lor stellar RV aud dey+6(0.004:v) for RV induced by the Earth's attnosphere. where σ(0.5) represents random uumbers followiug a gaussiandistribution witha mean of 0 aud a staucdard deviation of 6. | We have two sets of RV measurements, $\delta v_S+\sigma(0,\delta v_{rms,S})$ for stellar RV and $\delta
v_N+\sigma(0,\delta v_{rms,N})$ for RV induced by the Earth's atmosphere, where $\sigma(0,\delta)$ represents random numbers following a gaussiandistribution with a mean of 0 and a standard deviation of $\delta$. |
904,4 is the photon-lunited measurement error lor component £s aud v4.v is tlhe photou-limitecl measurement error for component Fay. | $\delta v_{rms,S}$ is the photon-limited measurement error for component $F_S$ and $\delta v_{rms,N}$ is the photon-limited measurement error for component $F_N$. |
We weigh the final RV measturement with the inverse square of photou-Iimited RV uncertainties of these two components. which is expressed by the following equation: In practical Doppler imeasurenients. dey consists two components. stellar RV aud. Earth's barvceutric RV. | We weigh the final RV measurement with the inverse square of photon-limited RV uncertainties of these two components, which is expressed by the following equation: In practical Doppler measurements, $\delta v_S$ consists two components, stellar RV and Earth's barycentric RV. |
Dependiug ou the position of the Earth iu its orbit. there is au offset between jeg aud dea. which is the Earth's baryceutrie velocity. | Depending on the position of the Earth in its orbit, there is an offset between $\delta v_S$ and $\delta v_N$, which is the Earth's barycentric velocity. |
The Earth's baryceutrie motion has a seimi-amplitude of 30 καννς17. | The Earth's barycentric motion has a semi-amplitude of 30 $\rm{km}\cdot\rm{s}^{-1}$. |
Statistically. observed star has an anuually-varyving RV with a of ou-average 21.21 kuvs 1. | Statistically, observed star has an annually-varying RV with a semi-amplitude of on-average 21.21 $\rm{km}\cdot\rm{s}^{-1}$ . |
We artificially shift a stellar spectrum by an amount less (han 21.21 kms.| in order to generate a offset between stellar spectrum and AA spectrum. | We artificially shift a stellar spectrum by an amount less than 21.21 $\rm{km}\cdot\rm{s}^{-1}$ in order to generate a offset between stellar spectrum and $AA$ spectrum. |
904,. and ρωN are then calculated for Fy aud Fa. | $\delta v_{rms,S}$ and $\delta v_{rms,N}$ are then calculated for $F_S$ and $F_N$. |
We choose the median of 6444v to represeut a typical ρε4 value from calculations based ou different input barycentrie velocities. | We choose the median of $\delta v_{rms,N}$ to represent a typical $\delta v_{rms,N}$ value from calculations based on different input barycentric velocities. |
We further assume that observed star has a constant RV (ie.. no differential RV). and £x has au RV fluctuation with au RMS of Sexy,pa; because of the Earth's turbulent. atiuospliere. | We further assume that observed star has a constant RV (i.e., no differential RV), and $F_N$ has an RV fluctuation with an RMS of $\delta v_{N,ATM}$ because of the Earth's turbulent atmosphere. |
. The measured RV uncertainty ov is equal to: Iu reality. RV uucertainty of £x is uot dominated by photou-nolse. iustead. it is dominated by atiuosphlieric behaviors such as wind. molecular columu- density change. | The measured RV uncertainty $\delta v$ is equal to: In reality, RV uncertainty of $F_N$ is not dominated by photon-noise, instead, it is dominated by atmospheric behaviors such as wind, molecular column density change, etc. |
ete. 2? used HARPS archive data and found that ος lines are stable toa 101n*s.© level over 6 years. | \citet{Figueira2010} used HARPS archive data and found that $O_2$ lines are stable to a 10 $\rm{m\cdot s}^{-1}$ level over 6 years. |
However. long term stability ol telluric lines (over years) becomes worse if we take into cousideration other gas molecules such as HoO aud COs. | However, long term stability of telluric lines (over years) becomes worse if we take into consideration other gas molecules such as $H_2O$ and $CO_2$. |
The uncertainty incdtcecd by atmospheric tellurie lines is trauslerred to δη via Equation (6)). | The uncertainty induced by atmospheric telluric lines is transferred to $\delta v_{rms}$ via Equation \ref{eq:simple_example_tulleric_2}) ). |
In order to calculate the final RV uncertainty. 0c,4,:. we need to calculate RV uncertainty 005,44 aud 9cjs according to Equation (3)). in which two terms need to be caleulated: Q and iN... | In order to calculate the final RV uncertainty, $\delta v_{rms}$ , we need to calculate photon-limited RV uncertainty $\delta v_{S,rms}$ and $\delta v_{N,rms}$ according to Equation \ref{eq:overall_Doppler}) ), in which two terms need to be calculated: Q and $N_{e^-}$. |
The spectral quality factors (Qs aud Qs) for the two components (Fy aud Fx) from equation (1)) are caleulated based on Equation (1)). | The spectral quality factors $Q_S$ and $Q_N$ ) for the two components $F_S$ and $F_N$ ) from equation \ref{eq:B_atm}) ) are calculated based on Equation \ref{eq:q_factor}) ). |
Ανs and V4. the photon flux of Ες aud Fy are calculated based oustellar type. maguituce. exposure time. iustrument specilications aud tellurie absorption properties. | $N_{e^-,S}$ and $N_{e^-,N}$ , the photon flux of $F_S$ and $F_N$ are calculated based onstellar type, magnitude, exposure time, instrument specifications and telluric absorption properties. |
Note the ratio of Ανy and Ανy remaius constant as long asatmospheric absorption stays uuchauged because telluric line absorption is imprinted ou the stellar spectrum. | Note the ratio of $N_{e^-,S}$ and $N_{e^-,N}$ remains constant as long asatmospheric absorption stays unchanged because telluric line absorption is imprinted on the stellar spectrum. |
aud The quartic typically has a single positive real root when the timestep satisfies the Courant condition of relsec:dt.. | and The quartic typically has a single positive real root when the timestep satisfies the Courant condition of \\ref{sec:dt}. |
Writing the coellicient of rv in equation (35)) as οι aid the coustant coefficient as cy. this single root lies between zero. aud the smaller of Jey/ey| and 1]. | Writing the coefficient of $x$ in equation \ref{eqn:interactionquartic}) ) as $c_1$ and the constant coefficient as $c_0$, this single root lies between zero, and the smaller of $|c_0/c_1|$ and $|c_0|^{1/4}$. |
The updated gas energy density is obtained by fiuding the root using Newtou-Raplson iteration ou this interval. with bisection when the Newtou-Raplsou methocl fails. | The updated gas energy density is obtained by finding the root using Newton-Raphson iteration on this interval, with bisection when the Newton-Raphson method fails. |
The updated radiation energy deusity is determined by substituting the updated gas energy density in equation (32)). | The updated radiation energy density is determined by substituting the updated gas energy density in equation \ref{eqn:erdifference}) ). |
When the absorption opacity & is zero aud the material-raciationu interaction ternis vanish. equations (32)) aud (33)) are no longer coupled. | When the absorption opacity $\kappa$ is zero and the material-radiation interaction terms vanish, equations \ref{eqn:erdifference}) ) and \ref{eqn:edifference}) ) are no longer coupled. |
In this iustauce. numerical energy conservation may be improved by using time-ceutered. values for the gas aud radiation pressures. | In this instance, numerical energy conservation may be improved by using time-centered values for the gas and radiation pressures. |
In. this differencing scheme. the time-advanced gas euergy ceusity € is with g=(5—L)(V*v)". | In this differencing scheme, the time-advanced gas energy density $e$ is with $q=(\gamma-1)(\nabla\cdot{\bf v})^n$. |
For the radiation euergy deusity update. q is replaced by the quantity mulliplvingNg E"1 in equation1 (31)). | For the radiation energy density update, $q$ is replaced by the quantity multiplying $E^{n+1}$ in equation \ref{eqn:gradvprad}) ). |
Advection of the radiation energy. proceeds exactly as described by SN for the bydrocyuamic variables. | Advection of the radiation energy proceeds exactly as described by SN for the hydrodynamic variables. |
Having expressed the advectiou terms in iutegral form in equation (23)). we apply the conservative differenciug scheme used for the livdrodsnamie variables. | Having expressed the advection terms in integral form in equation \ref{eqn:ertransport}) ), we apply the conservative differencing scheme used for the hydrodynamic variables. |
In this scheme. the flux of advected radiation energy is computed [rom the mass fux in order to reduce the relative numerical diffusion of radiation with respect to the eas. | In this scheme, the flux of advected radiation energy is computed from the mass flux in order to reduce the relative numerical diffusion of radiation with respect to the gas. |
The flux across every zoue interface is computed using an interpolation method which may be selected as either donor cell. vau Leer. or piecewise parabolic. | The flux across every zone interface is computed using an interpolation method which may be selected as either donor cell, van Leer, or piecewise parabolic. |
These fIuxes are then used to update the radiation euergy density ina directionally split fashion. | These fluxes are then used to update the radiation energy density in a directionally split fashion. |
variations as a function of the optical depth within cach of the models A. DB and €. For instance. in model A (disk case) its equivalent width varies remarkably depending on Tv. reaching à maximum value of ~500 at mo~1 and decreasing to ~65 at zy~6. | variations as a function of the optical depth within each of the models A, B and C. For instance, in model A (disk case) its equivalent width varies remarkably depending on $\tau_V$, reaching a maximum value of $\sim
500$ at $\tau_V\sim 1$ and decreasing to $\sim 65$ at $\tau_V\sim 6$. |
In addition. strong variations are also evident. for a fixed τι. as a function of the ecometry ofthe dust extinction. | In addition, strong variations are also evident, for a fixed $\tau_V$, as a function of the geometry of the dust extinction. |
For instance. for zy:=1. the 2200 feature displays laree dilferences among the models A. D and €. In this regard. it is important to note that. at least for the disk clistribution (model A). the feature is strongly weakened or suppressed because the emission from opaque galaxies is mostly contributed by stars outside the dust. disk and hence παοσο) by absorption. | For instance, for $\tau_V=1$, the 2200 feature displays large differences among the models A, B and C. In this regard, it is important to note that, at least for the disk distribution (model A), the feature is strongly weakened or suppressed because the emission from opaque galaxies is mostly contributed by stars outside the dust disk and hence unaffected by absorption. |
The strong. variations of the feature as a function of he geometric configuration suggest that the compensating ellect of scattering may. not be the only way to reduce he strength. of the feature. and that the geometry. may nave à dominant inlluence. | The strong variations of the feature as a function of the geometric configuration suggest that the compensating effect of scattering may not be the only way to reduce the strength of the feature, and that the geometry may have a dominant influence. |
Lt is important to recall here hat the main contribution to the bump comes from small (e< Q.lpm) graphite grains: for the cust composition here adopted. the weighted average radius at this wavelength is in fact ~0.05 yam. Correspondingly. the albedo is rather ow. (£o)&0.2. and the scattering phase function. rather isotropic. (ο)20.2. | It is important to recall here that the main contribution to the bump comes from small $a\leq 0.1 \mu$ m) graphite grains; for the dust composition here adopted, the weighted average radius at this wavelength is in fact $\sim 0.05~\mu$ m. Correspondingly, the albedo is rather low, $\langle \tilde\omega \rangle \simeq 0.2$, and the scattering phase function rather isotropic, $\langle g \rangle \simeq 0.2$. |
Thus. it seems unlikely that scattering is the main. or the only. origin of the weakness of the 2200 feature Compared to the pure absorption case. | Thus, it seems unlikely that scattering is the main, or the only, origin of the weakness of the 2200 feature compared to the pure absorption case. |
Llowever. a more detailed analysis is required to answer this question unambiguously. | However, a more detailed analysis is required to answer this question unambiguously. |
Although our results suggest that the strength of the 2200 feature may be a strong function of the ecometrical ellects. they co not rule out the possibility that external galaxies may have chemical grain properties cilferent from those of our Galaxy. | Although our results suggest that the strength of the 2200 feature may be a strong function of the geometrical effects, they do not rule out the possibility that external galaxies may have chemical grain properties different from those of our Galaxy. |
In fact. it is important to recall that the 2200 feature is weak in the LAIC. and absent in the SAIC (Prevot et al. | In fact, it is important to recall that the 2200 feature is weak in the LMC, and absent in the SMC (Prevot et al. |
1984). | 1984). |
The issue of the 2200 feature will be investigated in details in our forthcoming paper. | The issue of the 2200 feature will be investigated in details in our forthcoming paper. |
As a first application. we investigate the cllects of. dust extinction on the SEDs of two extremely red. galaxies. at igh-z. | As a first application, we investigate the effects of dust extinction on the SEDs of two extremely red galaxies at $z$. |
Extremely red. objects. EROs. (ie. £Aoc 6) are found. both in random sky fieles. or in λος fields (Elston et al. | Extremely red objects, EROs, (i.e. $R-K > 6$ ) are found both in random sky fields or in AGNs fields (Elston et al. |
LOSS: AleCarthy. Persson West 1992: Hu tideway 1994). | 1988; McCarthy, Persson West 1992; Hu Ridgway 1994). |
Since their colours may be strongly alfected ον dust extinction. these galaxies represent an interesting application of our moclels. and suitable objects to investigate at what level the SEDs of dusty vounger galaxies can minic hose of old. ellipticals once we treat dust extinction in a realistic wav. and to constrain the spatial distribution of the cust. | Since their colours may be strongly affected by dust extinction, these galaxies represent an interesting application of our models, and suitable objects to investigate at what level the SEDs of dusty younger galaxies can mimic those of old ellipticals once we treat dust extinction in a realistic way, and to constrain the spatial distribution of the dust. |
We have selected from the literature two extremely red galaxies. LIRIO and HIt14. for which the best optical-IR. photometric SEDs are available (lu Ridgeway 1994: Graham Dev 1996). | We have selected from the literature two extremely red galaxies, HR10 and HR14, for which the best optical-IR photometric SEDs are available (Hu Ridgway 1994; Graham Dey 1996). |
Our study. is more concentrated on LRLO. the only ERO with spectroscopically determined recshilt. | Our study is more concentrated on HR10, the only ERO with spectroscopically determined redshift. |
We attempted. to fit the SEDs of LRLO and 141 with svnthetic stellar population spectra. and applying the attenuation curves derived by our models in order to take into account the possible dust extinction occurring in these ealaxies. | We attempted to fit the SEDs of HR10 and HR14 with synthetic stellar population spectra, and applying the attenuation curves derived by our models in order to take into account the possible dust extinction occurring in these galaxies. |
In our experiments we made use of the synthetic. unextineted. stellar population spectra of Bruzual Charlo (1993). with instantaneous burst of star formation. solar metallicity. ancl Salpeter LME. | In our experiments we made use of the synthetic, unextincted, stellar population spectra of Bruzual Charlot (1993), with instantaneous burst of star formation, solar metallicity, and Salpeter IMF. |
Figure 2 shows our results for HIUIO. a galaxy at 21.44 with a colour index £0A—6.5. | Figure 2 shows our results for HR10, a galaxy at $z$ =1.44 with a colour index $I-K^{'}$ =6.5. |
Since its quoted.B-bau Hux is not statistically significant (2.2( level). we adop a 37 upper limit. | Since its quoted$B$ -band flux is not statistically significant $\sim2.2\sigma$ level), we adopt a $3\sigma$ upper limit. |
We also correct for the ~20% L-band Dux contamination [rom the IIo. line. | We also correct for the $\sim$ H-band flux contamination from the $\alpha$ line. |
Graham Dev (1996). acopting a foreground screen dust absorption mocel. fou that the SED of 1180 cannot be reproduced: without dus extinction. | Graham Dey (1996), adopting a foreground screen dust absorption model, found that the SED of HR10 cannot be reproduced without dust extinction. |
Our independent. analvsis confirms that. dus extinction must play an important role in HIUIO. and even very old. (ie. very red) stellar. populations (up to 7 Cr) cannot reproduce its SED without extinction. | Our independent analysis confirms that dust extinction must play an important role in HR10, and even very old (i.e. very red) stellar populations (up to 7 Gyr) cannot reproduce its SED without extinction. |
However our modeling allows us to £o a step further. ancl to derive clucs on the spatial distribution of the dust in LIRTO. | However our modeling allows us to go a step further, and to derive clues on the spatial distribution of the dust in HR10. |
In fact. we find that model A (dust disk) cannot fit the data because the observed. colours are always too red compared to those produced by this model for any τι age. sy. ancl 7. | In fact, we find that model A (dust disk) cannot fit the data because the observed colours are always too red compared to those produced by this model for any $\tau_{V}$, age, $z_{d}$, and $i$. |
Model € allows a better fit to the SED. but the reddening is still not sullicient to match the £A7 colour. | Model C allows a better fit to the SED, but the reddening is still not sufficient to match the $I-K^{'}$ colour. |
We find that only model 13 gives an acceptable fit to the SED (age ~ 1 Gyr and 7-6). | We find that only model B gives an acceptable fit to the SED (age $\sim$ 1 Gyr and $\tau_{V}$ =6). |
The SED can be fit also bv a 1.0 Cyr old population reddened by a dust. foreground screen. with Typcml.54 (or [£g (-0.55. adopting the Calactic relation decMg s). | The SED can be fit also by a 1.0 Gyr old population reddened by a dust foreground screen with $\tau_V$ =1.54 (or $E_{B-V}$ =0.55, adopting the Galactic relation $A_V \sim 3 E_{B-V}$ ). |
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