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There is clearly some amount of degeneracy between completeness. fi. 0.. and c-.
There is clearly some amount of degeneracy between completeness, $f_{3\sigma}$, $\delta_z$ , and $\sigma_z$.
The plots in Fig.
The plots in Fig.
5 and 6 provide a more complete view of the performance.
\ref{fig:zz_BCH_C5} and \ref{fig:zz_BCH_B5} provide a more complete view of the performance.
Remarkably. the negative biases introduced by and as reported above are much smaller or negligible for the COMBO code.
Remarkably, the negative biases introduced by and as reported above are much smaller or negligible for the COMBO code.
are tested in their release papers (??) only against real data from the Hubble Deep Field. besides simulations.
are tested in their release papers \citep{2000ApJ...536..571B, 2000A&A...363..476B} only against real data from the Hubble Deep Field, besides simulations.
now incorporates a new template set (see Sect. 3.3))
now incorporates a new template set (see Sect. \ref{sec:bpz}) )
that was specially calibrated for HST photometry.
that was specially calibrated for HST photometry.
The COMBO code. however. was originally designed for the ground-based survey CADIS (?).. where colours were measured bias-free from seeing adaptive photometry. and included photo-z's for point-source QSOs.
The COMBO code, however, was originally designed for the ground-based survey CADIS \citep{2001A&A...365..660W}, where colours were measured bias-free from seeing adaptive photometry, and included $z$ 's for point-source QSOs.
One of the biggest differences between the codes is the template set chosen and one might presume that most of the difference in performance originates from this point.
One of the biggest differences between the codes is the template set chosen and one might presume that most of the difference in performance originates from this point.
However. we run with the PEGASE templates used by the COMBO code às well as with the CWW templates plus two Kinney starburst templates originally used by in ?)..
However, we run with the PEGASE templates used by the COMBO code as well as with the CWW templates plus two Kinney starburst templates originally used by in \cite{2000ApJ...536..571B}.
We switch off the internal reddening because it is already included in the templates and the PEGASE agexextinction grid used by the COMBO code.
We switch off the internal reddening because it is already included in the templates and the PEGASE $\times$ extinction grid used by the COMBO code.
The results can neither compete with the best setups incorporating the BC templates nor with the COMBO code plus PEGASE templates.
The results can neither compete with the best setups incorporating the BC templates nor with the COMBO code plus PEGASE templates.
Hence. the implementation. of user-defined templates appears to be not straightforward and results may not be competitive with the template sets that are shipped with the code and were tested and optimised by the author.
Hence, the implementation of user-defined templates appears to be not straightforward and results may not be competitive with the template sets that are shipped with the code and were tested and optimised by the author.
Another interesting point is the comparison of the CC17 setup with the CBS setup.
Another interesting point is the comparison of the CC17 setup with the CB5 setup.
While the total exposure time with WEI is lower for CCI7. the performance of CCI7 is better in all statistics described here.
While the total exposure time with WFI is lower for CC17, the performance of CC17 is better in all statistics described here.
It is clear. that for the particular application of photo-Zs for bright objects. the exposure time was well spent on more filters (which ts an important result for future surveys).
It is clear, that for the particular application of $z$ 's for bright objects, the exposure time was well spent on more filters (which is an important result for future surveys).
However. the GaBoDS data of the CDFS are completely based on archive data and no specific observing programme was proposed to create these deep images.
However, the GaBoDS data of the CDFS are completely based on archive data and no specific observing programme was proposed to create these deep images.
Furthermore. for deeper applications. such as Lyman-break galaxy studies. where you simply need a very deep colour index between three bands. the GaBoDS data are certainly highly superior to the COMBO data.
Furthermore, for deeper applications, such as Lyman-break galaxy studies, where you simply need a very deep colour index between three bands, the GaBoDS data are certainly highly superior to the COMBO data.
Table + summarises the results on the FDF and Fig 7. shows photometric vs. spectroscopic redshift for selected setups.
Table \ref{tab:res_FDF} summarises the results on the FDF and Fig \ref{fig:FDF} shows photometric vs. spectroscopic redshift for selected setups.
In the lower redshift bin again the COMBO code combined with imaging data in 8 filters delivers the smallest outher rate. bias. and scatter when compared to and in 8 filters.
In the lower redshift bin again the COMBO code combined with imaging data in 8 filters delivers the smallest outlier rate, bias, and scatter when compared to and in 8 filters.
At least in this redshift interval the results are nearly as goodas the results produced by ?) with atemplate set specifically calibrated for the FDF.
At least in this redshift interval the results are nearly as goodas the results produced by \cite{2004A&A...421...41G} with atemplate set specifically calibrated for the FDF.
against the possibility of à much louger loop D. uamely L=2.5 «1079 απ.
against the possibility of a much longer loop B, namely L=2.5 $\times 10^{10}$ cm.
Indeed. loop B. could be even shorter hau loop A. if a residual heating[m] were preseut duriusc» xiase D1.
Indeed, loop B could be even shorter than loop A, if a residual heating were present during phase D1.
However. this would imply two different regimes of residual heating. one in phase D1 aud another in phase D2. aud introduce another set of free parameters in the uodeling. which we prefer uot to do. if unnecessary.
However, this would imply two different regimes of residual heating, one in phase D1 and another in phase D2, and introduce another set of free parameters in the modeling, which we prefer not to do, if unnecessary.
Since the slower evolution expected from a longer loop nav naturally lead to a delaved emission peak. we have explored the possibility that the flare in this long loop D is triggered at the same time as that in loop A. with a lower iuteusitv and longer duration.
Since the slower evolution expected from a longer loop may naturally lead to a delayed emission peak, we have explored the possibility that the flare in this long loop B is triggered at the same time as that in loop A, with a lower intensity and longer duration.
In this hypothesis. the slower decay D2 may be explained with the superposition of the coutiuuatiou of the fast decay DI with the rise of the flare of loop D (hypothesis (1) in Sec. 1.2)).
In this hypothesis, the slower decay D2 may be explained with the superposition of the continuation of the fast decay D1 with the rise of the flare of loop B (hypothesis (ii) in Sec. \ref{sec:d2}) ).
Fig.5 shows the results obtained with a heating pulse located at the top of loop D with My=1 cre ? +. constaut for 3200 s. The figure shows the light curve of the wo flare compoucuts separately and the ieht curve obtained bv sunnuiug the spectra of loop A and loop D (with distiuct cross-section areas) at correspoudine times. compared to the observed light curve.
\ref{fig:lntp} shows the results obtained with a heating pulse located at the top of loop B with $H_0 = 1$ erg $^{-3}$ $^{-1}$, constant for 3200 s. The figure shows the light curve of the two flare components separately, and the light curve obtained by summing the spectra of loop A and loop B (with distinct cross-section areas) at corresponding times, compared to the observed light curve.
The ight curve of loop D rises very eradually and peaks at nne foc3500 κ. nore than 2000 s later than the flare in] loop A. The oop cross-section of this second loop that bes fits the ieht curve is 6.2«LOS cut. which corresponds to quite a small aspect ratio R/Lzx0.06.
The light curve of loop B rises very gradually and peaks at time $t \approx 3500$ s, more than 2000 s later than the flare in loop A. The loop cross-section of this second loop that best fits the light curve is $6.2 \times 10^{18}$ $^2$ , which corresponds to quite a small aspect ratio $R/L \approx 0.06$.
Fitting the otal spectra ο oop ÁAlloop D with isothermal models. the elobal evolution of the cussion nieasure is well described.
Fitting the total spectra of loop $+$ loop B with isothermal models, the global evolution of the emission measure is well described.
The evohitici of the temperature instead shows a deep munimuun a ine fzc2000. s. This is not present in the data. aux sugeests us to reject this model.
The evolution of the temperature instead shows a deep minimum at time $t \approx 2000$ s. This is not present in the data, and suggests us to reject this model.
A similar temperature dip appears also if such long loop D is heated at the footpoiuts.
A similar temperature dip appears also if such long loop B is heated at the footpoints.
Fie.
Fig.
6 shows results obtained with a loop B twin of loop A. and with a heating duration aud spatial distribution of loop D ideutical to that of loop A. but trigecred 2600 s later. with a rate teu times lower. fy=6 ere 5 .
\ref{fig:bs6} shows results obtained with a loop B twin of loop A, and with a heating duration and spatial distribution of loop B identical to that of loop A, but triggered 2600 s later, with a rate ten times lower, $H_0 = 6$ erg $^{-3}$ $^{-1}$.
Iu. this alternative scenario. the decaving heating of loop A is maintained.
In this alternative scenario, the decaying heating of loop A is maintained.
The best-fitting section area of loop D is 16.7«1015 αμ”, five times lareer than the area of the first flaring loop. or. equivalently. an arcade of five loops equal to loop A. This combination describes better the temperature treud. but the helt curve shows a small dip at time f~3000 s. As suggested by the residuals. the dip cau be filled and the fit further iuprove siuuply by adding a third munor flaring component. adjusting the loop cross-section areas aud the heating time shifts appropriately (Fig. 7)).
The best-fitting cross-section area of loop B is $16.7 \times 10^{18}$ $^2$, five times larger than the area of the first flaring loop, or, equivalently, an arcade of five loops equal to loop A. This combination describes better the temperature trend, but the light curve shows a small dip at time $t \sim 3000$ s. As suggested by the residuals, the dip can be filled – and the fit further improved – simply by adding a third minor flaring component, adjusting the loop cross-section areas and the heating time shifts appropriately (Fig. \ref{fig:best}) ).
The best conmnibiuatiou that we fud is to add two oops D with area 9.0«1015 ci?2 and ,10,8∖&1015N cni Doo,(2.5 and 3. times. the area of loop A. still compatible with an arcade of ~5 loops equal to loop A). and heated with time shifts of 2200 s and 2800 sx. respectively. since the start of the eating of loop A. Tn alternative. we nav think to fill the eap of the liebt curve by considering a longer-astiug heating pulse iu loop D. but we checked that this choice fails to reproduce adequately the temperature evolution.
The best combination that we find is to add two loops B with area $9.0 \times 10^{18}$ $^2$ and $10.8 \times 10^{18}$ $^2$ (2.5 and 3 times the area of loop A, still compatible with an arcade of $\sim 5$ loops equal to loop A), and heated with time shifts of 2200 s and 2800 s, respectively, since the start of the heating of loop A. In alternative, we may think to fill the gap of the light curve by considering a longer-lasting heating pulse in loop B, but we checked that this choice fails to reproduce adequately the temperature evolution.
The latest decay DI can be reasonably well fitted by assundue a residual heating of loop D (or of two loops D) with the same characteristics aude-fokdius time as the one of loop A. and the superposition of the decays of loops A and B.
The latest decay D4 can be reasonably well fitted by assuming a residual heating of loop B (or of two loops B) with the same characteristics ande-folding time as the one of loop A, and the superposition of the decays of loops A and B.
Sunspots appear at the solar surface following a 11-year cycle.
Sunspots appear at the solar surface following a 11-year cycle.
They reveal the presence of strong magnetic fields in the solar interior. and suggest the existence of à dynamo process governing its evolution.
They reveal the presence of strong magnetic fields in the solar interior, and suggest the existence of a dynamo process governing its evolution.
However. the process by which sunspots are formed is. still unkown.
However, the process by which sunspots are formed is still unknown.
At the solar surface. sunspots are observed as bipolar patches of radial magnetic field.
At the solar surface, sunspots are observed as bipolar patches of radial magnetic field.
This intuitively suggest that they are formed by the emergence of horizontal concentratios of magnetic field lines. often called magnetic flux tubes.
This intuitively suggest that they are formed by the emergence of horizontal concentrations of magnetic field lines, often called magnetic flux tubes.
Sice Parker's original proposal of magnetic buoyancy ?(?).. the model has evolved through the thin flux tube approximation (?).. to numerical simulations of the emergence of 3D magetic flux tubes (?)..
Since Parker's original proposal of magnetic buoyancy \citep{P55}, the model has evolved through the thin flux tube approximation \citep{Spruit_81}, to numerical simulations of the emergence of 3D magnetic flux tubes \citep{Fan_08}.
The buoyant rise of flux tubes from the tachocline up to the surface 1s currently the most widely accepted mechanism of sunspot formation.
The buoyant rise of flux tubes from the tachocline up to the surface is currently the most widely accepted mechanism of sunspot formation.
During the last three decades much work has been done in order to understand the buoyancy phenomena and to reconcile the results of flux tube models with phenomenological sunspots rules such as the Joy's law or the topological difference between the two spots in a pair.
During the last three decades much work has been done in order to understand the buoyancy phenomena and to reconcile the results of flux tube models with phenomenological sunspots rules such as the Joy's law or the topological difference between the two spots in a pair.
In spite of the fact that these basic observations have been qualitatively reproduced by the flux tube models. a set of complications have questioned the feasibility of this scenario.
In spite of the fact that these basic observations have been qualitatively reproduced by the flux tube models, a set of complications have questioned the feasibility of this scenario.
These may be summarized as follows:
These may be summarized as follows:
F.G.B. thanks Daniele Pierini for helpful discussions and suggestions.
F.G.B. thanks Daniele Pierini for helpful discussions and suggestions.
We acknowledge the support of NASA through grant numbers NAGS-12785. NAG5-13301. and NNGO-6GII1G. the NSF Office of Polar Programs. the Canadian Space Ageney. the Natural Sciences and Engineering Research Council (NSERC) of Canada. and the UK Science and Technology Facilities Council (STFC).
We acknowledge the support of NASA through grant numbers NAG5-12785, NAG5-13301, and NNGO-6GI11G, the NSF Office of Polar Programs, the Canadian Space Agency, the Natural Sciences and Engineering Research Council (NSERC) of Canada, and the UK Science and Technology Facilities Council (STFC).
This research has been enabled by the use of WestGrid computing resources.
This research has been enabled by the use of WestGrid computing resources.
This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Administration.
This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.
This publication makes use of data products from the Two Micron All Sky Survey. which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology. funded by the National Aeronautics and Space Administration and the National Science Foundation,
This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation.
This work is based in part on observations made with theSpitzer Space Telescope. which is operated by the Jet Propulsion Laboratory. California Institute of Technology under a contract with NASA.
This work is based in part on observations made with the Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA.
Here we present the full catalogue of all BLAST sources detected in the central 0.8 deg? of the A3112 field at 250. 350. ormieron.. with a significance of at least 3o.
Here we present the full catalogue of all BLAST sources detected in the central 0.8 $^2$ of the A3112 field at 250, 350, or, with a significance of at least $3\sigma$.
For each source. its position. flux density and errors are provided.
For each source, its position, flux density and errors are provided.
Where flux densities are provided in more than one band. the quoted position is the averaged position of the matched sources (cf.
Where flux densities are provided in more than one band, the quoted position is the averaged position of the matched sources (cf.
Section 2.1).
Section \ref{blastdata}) ).
An asterisk C) following the BLAST ID means the source is identitied as a counterpart of a cluster member.
An asterisk $^{*}$ ) following the BLAST ID means the source is identified as a counterpart of a cluster member.
2231. have demoustrated that the bulls of the OIL maser enuission arises in circíunnuclear disks or tori (Pillstrometal.2001:IKlóckner 2003).
231, have demonstrated that the bulk of the OH maser emission arises in circumnuclear disks or tori \citep{pihlstrom01,kloeckner03}.
. Tlowever. there are also indications that a single disk component cauuo account for mascr enission. iin 2273: (IXlócknereal.2001).
However, there are also indications that a single disk component cannot account for maser emission, in 273; \citep{kloeckner04}.
. Early single dish studies suggested the presence of που]τος wines iu several ΟΠ meceamascr spectra (Baan.Haschick.auIlenkel1989).
Early single dish studies suggested the presence of blueshifted wings in several OH megamaser spectra \citep{baanetal89}.
. These have beeu interprete as outflows. which possibly could provide au explanation for the emission that does not ft comfortably within a disk model.
These have been interpreted as outflows, which possibly could provide an explanation for the emission that does not fit comfortably within a disk model.
The existence of outflows as well as inflows is not surprise.e. eelven hat ΟΠ meeamascr galaxies are exclusively associated with (Ultra)Luuiuous Iufra-Red Galaxies ((UJLIRGs). (
The existence of outflows as well as inflows is not surprising, given that OH megamaser galaxies are exclusively associated with (Ultra)Luminous Infra-Red Galaxies ([U]LIRGs). (
U)LIRGSs probably represent short periocs of unclear starburst activity trigeered by iuerger events.
U)LIRGs probably represent short periods of nuclear starburst activity triggered by merger events.
Such ierecrs will rapidly transport eas to the central regions. causing eas to fall awards towards the nuclei.
Such mergers will rapidly transport gas to the central regions, causing gas to fall inwards towards the nuclei.
The combined effects of supernova explosions aud stellar winds eeuerated iu this nuclear starburst can eutrain the iuterstellu: mediunu iu outflows with velocities of LOO)1000 citeplieckiian90. altou99..
The combined effects of supernova explosions and stellar winds generated in this nuclear starburst can entrain the interstellar medium in outflows with velocities of $100-1000$ \\citep{heckman90, alton99}.
It is also possible that OIL ανα cussion occurs in both of the imereine nuclei. as has been seen in Àrp2220 (Diamondctal.1989).
It is also possible that OH maser emission occurs in both of the merging nuclei, as has been seen in 220 \citep{diamond89}.
. With a slight. offse In systemic velocity between the two imeremeg nuclei. the combined spectra. observed with a single dish telescope would be relatively broad.
With a slight offset in systemic velocity between the two merging nuclei, the combined spectrum observed with a single dish telescope would be relatively broad.
However. for the very broadest OU meeamascr lines (exceeding 1000 J). at is hard to interpret the lines either as originating from a single disk component or roni a pair of nuclei.
However, for the very broadest OH megamaser lines (exceeding 1000 ), it is hard to interpret the lines either as originating from a single disk component or from a pair of nuclei.
A combination of orbital nechanics. possible disk rotation. molecular inflow and outfow. aud the velocity separation of the two nain lines at 1667 MIIz and 1665 MITz is likely to nake up the velocity range in such ΟΠ megamaser spectra.
A combination of orbital mechanics, possible disk rotation, molecular inflow and outflow, and the velocity separation of the two main lines at 1667 MHz and 1665 MHz is likely to make up the velocity range in such OH megamaser spectra.
Earlier global VLBI experiments have reported je presence of LOO broad ΟΠ maser lines on parsec scales (Diaunondal. 1999)..
Earlier global VLBI experiments have reported the presence of 100 broad OH maser lines on parsec scales \citep{diamond99}.
In this paper we will concentrate on investigating the cause of ΟΠ meeamaser lines with widths exceeding 1000Ἐ
In this paper we will concentrate on investigating the cause of OH megamaser lines with widths exceeding 1000.
ν, We report on VLBA observatious of two OTT meeamaser ealaxies. 111070|0525 and 112032|1707. that have full width zero intensity velocities of 15002000 laus.
We report on VLBA observations of two OH megamaser galaxies, 14070+0525 and 12032+1707, that have full width zero intensity velocities of $1500-2000$ .
These objects are additionally iuterestiug because of their high redshifts (+=0.217 aud :=0.265).
These objects are additionally interesting because of their high redshifts $z=0.217$ and $z=0.265$ ).
Furthermore. their high ΟΠ luminosities of Low=1.3«WAL. aud Lou=1.2<1011. respectively make the term “eleamascr suitable. and as such these galaxies are two of the most powerful OIL megamascr galaxies known.
Furthermore, their high OH luminosities of $L_{\rm OH}=1.3\times 10^4L_{\odot}$ and $L_{\rm OH}=1.2\times 10^4L_{\odot}$ respectively make the term 'gigamaser' suitable, and as such these galaxies are two of the most powerful OH megamaser galaxies known.
The main anu of the current observations was to determine whether the broad lines are associated with inflows. nmtüows. rotating structures or violent kiucimatics.
The main aim of the current observations was to determine whether the broad lines are associated with inflows, outflows, rotating structures or violent kinematics.
This paper also presets Arecibo absorption data for one of the sources. 112032|1707.
This paper also presents Arecibo absorption data for one of the sources, 12032+1707.
112032|1707 was observed in pliase-refercucing mode with the VLBA for 12 hours ou 2002 July 8.
12032+1707 was observed in phase-referencing mode with the VLBA for 12 hours on 2002 July 8.
The redshift of 112032|1707 (2= 0.217) shifted the MMIIz line to NMIIz. that was used as the center of the observing band.
The redshift of 12032+1707 $z=0.217$ ) shifted the MHz line to MHz, that was used as the center of the observing band.
To cover the complete OT chussion velocity range of 112032|1707 (~2000 1). a bandwidth of MMIIz per left alc vielt hand polarization was used.
To cover the complete OH emission velocity range of 12032+1707 $\sim$ 2000 ), a bandwidth of MHz per left and right hand polarization was used.
telescopes were available at the time of the observaions. aud only iinor periods of tie required fiaeeimg due to radio frequency interfterenc (RFI).
All telescopes were available at the time of the observations, and only minor periods of time required flagging due to radio frequency interference (RFI).
LLL070|0525 was observed with a similar setup in June 9. 2002.
14070+0525 was observed with a similar setup in June 9, 2002.
The MMIIZ IF pair was centered on the redsbhifted ἐν= 0.265) requeney of MMIIz.
The MHz IF pair was centered on the redshifted $z=0.265$ ) frequency of MHz.
Due to RFI the KP clescope showed extreme and highly variable system temperatures. aud to a large extent data iad to be flageed for this telescope.
Due to RFI the KP telescope showed extreme and highly variable system temperatures, and to a large extent data had to be flagged for this telescope.
Furthermore. he LA auteuna was broken aud so did not uticipate in these observations.
Furthermore, the LA antenna was broken and so did not participate in these observations.
At the observing requeucy of MATTz. several REI spikes could veo seen iu the autocorrelatiou spectra.
At the observing frequency of MHz, several RFI spikes could be seen in the autocorrelation spectra.
Since he frequencies at which these spikes occurred differed between the sites. they did uotaffect the cross correlation spectra for the most part.
Since the frequencies at which these spikes occurred differed between the sites, they did notaffect the cross correlation spectra for the most part.
Dadly affeced time ranges were removed on by baseline
Badly affected time ranges were removed on `by baseline'
away towards the observer.
away towards the observer.
The atmosphere is then in radiative equilibrium.
The atmosphere is then in radiative equilibrium.
We now present our first theoretical SN Ia light curves obtained with our extensions to the general purpose model atmosphere code aand compare them to observed SN la light curves.
We now present our first theoretical SN Ia light curves obtained with our extensions to the general purpose model atmosphere code and compare them to observed SN Ia light curves.
The online supernova spectrum archive (SUSPECT) (??) provides numerous of observations of different types of supernovae.
The online supernova spectrum archive (SUSPECT) \citep{suspect01,suspect02} provides numerous of observations of different types of supernovae.
For this work. the observed light curves of SN 2002bo and SN 1999ee are used to compare them to our results of model light curves.
For this work, the observed light curves of SN 2002bo and SN 1999ee are used to compare them to our results of model light curves.
Photometric light curve observations of SN 2002bo in different photometric bands (?) have been obtained.
Photometric light curve observations of SN 2002bo in different photometric bands \citep{benetti04} have been obtained.
SN 1999ee also has observed spectra (?) and photometry (?)..
SN 1999ee also has observed spectra \citep{hamuy99ee99ex02} and photometry \citep{stritz99ee99ex02}.
The energy solver is now applied to calculate synthetic light curves of SNe la. The SN Ia light curve evolution ts calculated during the free expansion phase.
The energy solver is now applied to calculate synthetic light curves of SNe Ia. The SN Ia light curve evolution is calculated during the free expansion phase.
For the initial model atmosphere structure. the results of the explosion calculation of other groups are used as the input structure.
For the initial model atmosphere structure, the results of the explosion calculation of other groups are used as the input structure.
Each layer has a certain expansion velocity. which does not change during the evolution. because homologous expansion is assumed.
Each layer has a certain expansion velocity, which does not change during the evolution, because homologous expansion is assumed.
We start the model light curve calculation a few days after explosion.
We start the model light curve calculation a few days after explosion.
In the first few days the SN la envelope is optically thick and compact.
In the first few days the SN Ia envelope is optically thick and compact.