ReadingTimeMachine/rtm-sgt-ocr-v1 · Datasets at Hugging Face

source stringlengths 1 2.05k ⌀	target stringlengths 1 11.7k
. The drawback of these approaches is that the ruilcing of tailored. model atmospheres can be very. time-consuming and needs much more computing power.	The drawback of these approaches is that the building of tailored model atmospheres can be very time-consuming and needs much more computing power.
Also. in order to compute reliable models. the input observed spectra need to have sulliciently high S/N and spectral resolution.	Also, in order to compute reliable models, the input observed spectra need to have sufficiently high S/N and spectral resolution.
The more recently developed: Standardized Canelle AMethocl (SCAL) (Llamuy&Pinto2002:Lamuyοἱal.2003) relies on the empirical correlationbetween the measured expansion velocity and the luminosity in optical (Nugentetal.2006:Poznanski2009). or in NI. (Maguireetal.2010a)— bands in the micelle of plateau phase. αἲ [50 days after explosion.	The more recently developed Standardized Candle Method (SCM) \citep{hamuy_pinto_scm, hamuy_scm} relies on the empirical correlationbetween the measured expansion velocity and the luminosity in optical \citep{nugent06, poznanski09} or in NIR \citep{maguirescm} bands in the middle of plateau phase, at +50 days after explosion.
“Phis method. requires. less extensive input. data. but needs a larger sample of Tvpe LL SNe with independently known distances in order to calibrate the empirical correlation.	This method requires less extensive input data, but needs a larger sample of Type II SNe with independently known distances in order to calibrate the empirical correlation.
Because this is basically a wxhotometrice method which compares apparent and absolute magnitudes. SCAL is more sensitive to interstellar reddening han EDPM/SIEXM. as mentioned above.	Because this is basically a photometric method which compares apparent and absolute magnitudes, SCM is more sensitive to interstellar reddening than EPM/SEAM, as mentioned above.
Beside photometry. both IEPMSIZAM and SCM need information on the expansion velocity at the photosphere (Pye).	Beside photometry, both EPM/SEAM and SCM need information on the expansion velocity at the photosphere $v_{phot}$ ).
At first. EPM seems to be more challenging. because it requires multi-epoch observations. while SCAL needs only he velocity at a certain epoch. at. \|50 days after explosion (eso).	At first, EPM seems to be more challenging, because it requires multi-epoch observations, while SCM needs only the velocity at a certain epoch, at +50 days after explosion $v_{50}$ ).
However. direct measurement of eso is possible only in he case of precise timing of the spectroscopic observation. which is rarely achievable.	However, direct measurement of $v_{50}$ is possible only in the case of precise timing of the spectroscopic observation, which is rarely achievable.
The correct. determination of Pii; is not trivial.	The correct determination of $v_{phot}$ is not trivial.
As SNe expand homologously (er. where e is the velocity of a given laver and r is its distance from the center). in most cases it ds dillicult to derive a unique velocity [rom the observed. spectral features.	As SNe expand homologously $v \sim r$, where $v$ is the velocity of a given layer and $r$ is its distance from the center), in most cases it is difficult to derive a unique velocity from the observed spectral features.
The most. [frequently Followed approach relies on measuring the Doppler-shift of the absorption minimum of certain spectral features (mostly A5169 or 11:2).	The most frequently followed approach relies on measuring the Doppler-shift of the absorption minimum of certain spectral features (mostly $\lambda 5169$ or $\beta$ ).
Another possibility is the computation of the eross-correlation between the SN spectrum and a set of spectral templates. consisting of either observed or model spectra. with known velocities.	Another possibility is the computation of the cross-correlation between the SN spectrum and a set of spectral templates, consisting of either observed or model spectra, with known velocities.
Phe third method is building a [Full NUTE model (like PILOENIN or CAIPCIEN) for à given SN spectrum. and. adopting the theoretical ρω from the rest-Litting moclel.	The third method is building a full NLTE model (like PHOENIX or CMFGEN) for a given SN spectrum, and adopting the theoretical $v_{phot}$ from the best-fitting model.
The aim of this paper is to present a similar. but ess computation-demanding approach to assign velocities ο observed SN spectra.	The aim of this paper is to present a similar, but less computation-demanding approach to assign velocities to observed SN spectra.
We apply the simple parametrized code (Fisher1999:Llatanoctal.1999). to model the observed spectra with an approximative. but self-consistent reatment of the formation of lines in the extended. 10mologously. expanding SN atmosphere.	We apply the simple parametrized code \citep{fisher99, hatano99} to model the observed spectra with an approximative, but self-consistent treatment of the formation of lines in the extended, homologously expanding SN atmosphere.
To illustrate the applicability ofSYNOW. we construct ancl fit. parametrized models to a sample of five wellcobserved. nearby Type H-P SNe that have a series of high S/N spectra publicly available.	To illustrate the applicability of, we construct and fit parametrized models to a sample of five well-observed, nearby Type II-P SNe that have a series of high S/N spectra publicly available.
We also check ancl re-calibrate some empirical correlations between velocities derived. from various methods.	We also check and re-calibrate some empirical correlations between velocities derived from various methods.
The deseription of the observational sample is given in δεο,	The description of the observational sample is given in \ref{sec_data}.
In §?? we first review the cdillerent. velocity measurement techniques for SNe U-P. then we present the details of the application of SYNOW models (§??7)).	In \ref{sec_issue} we first review the different velocity measurement techniques for SNe II-P, then we present the details of the application of models \ref{sec_synow}) ).
Phe results are collected ancl discussed in $?? and 8?7?.. while the implications of the results for the clistance measurements are in $??..	The results are collected and discussed in \ref{sec_results} and \ref{sec_disc}, while the implications of the results for the distance measurements are in \ref{sec_dist}.
We draw our conclusions in 8?7..	We draw our conclusions in \ref{sec_concl}.
Our sample contains S1 plateau phasespectra of five objects. SNe 1999em. 2004dj. 2004et. 2005es and 2006bp. respectively.	Our sample contains 81 plateau phase–spectra of five objects, SNe 1999em, 2004dj, 2004et, 2005cs and 2006bp, respectively.
All five SNe are wellobserved objects. they have been stucied in detail. and show a wide variety in their physical properties (see Table 1)).	All five SNe are well-observed objects, they have been studied in detail, and show a wide variety in their physical properties (see Table \ref{physdata}) ).
SN 1999em was discovered. on 29th October 1999. by Li(19900) in NGC 1637 at a very carly phase.	SN 1999em was discovered on 29th October 1999 by \citet{discov99em} in NGC 1637 at a very early phase.
Ht is a very. well-observed. well-studied object The explosion date was determined as 2451477.0€2 JD (Leonardοἱal.2002:Παινetal.2001:Dessart&LHillier 2006)...	It is a very well-observed, well-studied object The explosion date was determined as $2451477.0 \pm 2$ JD \citep{leonard99em, hamuy99em, dessart2006}. .
We used the spectra of Leonardetal.(2002). and. Hamus.ctal.(2001) downloaded. from the database.	We used the spectra of \citet{leonard99em} and \citet{hamuy99em} downloaded from the database.
Our. saniple contains 22 spectra. covering the first SO days of the plateu phase.	Our sample contains 22 spectra, covering the first 80 days of the plateu phase.
SN 2004dj was discovered on 31st July 2004 by Iagaki (Nakanoetal.2004) in à voung. massive cluster Sancdage-96 of NGC 2403. about 1: month after explosion.	SN 2004dj was discovered on 31st July 2004 by Itagaki \citep{discov04dj} in a young, massive cluster Sandage-96 of NGC 2403, about 1 month after explosion.
We used the spectra taken by Vinkóetal.(2006)..	We used the spectra taken by \citet{vinko04dj}.
Due to the lack of observed spectrophotometric standards. the lux-calibration of those spectra was inferior. but this is not a major concern when only velocities are to be determined.	Due to the lack of observed spectrophotometric standards, the flux-calibration of those spectra was inferior, but this is not a major concern when only velocities are to be determined.
We included. 12 spectra taken between \|47 ancl \|100 days after explosion in our sample.	We included 12 spectra taken between +47 and +100 days after explosion in our sample.
SN 2004et was discovered by Moretti (Zwitter.Murani&Moretti2004) on 27th September 2004 in NGC 6946.	SN 2004et was discovered by Moretti \citep*{discov04et} on 27th September 2004 in NGC 6946.
The spectra of Sahuetal.(2006). (downloaded. (rom SUSPECT) and Maguireetal.(2010b) were used. together. with 6 previously unpublished earlv-phase spectra taken with the l.8S-m telescope at DDO (see the Appendix ??7)).	The spectra of \citet{sahu04et} (downloaded from ) and \citet{maguire04et} were used, together with 6 previously unpublished early-phase spectra taken with the 1.88-m telescope at DDO (see the Appendix \ref{appendix_a}) ).
The 22 spectra cover the period of \| 11 to \| 104 days after explosion.	The 22 spectra cover the period of $+$ 11 to $+$ 104 days after explosion.
SN 2005cs was discovered on 29th June 2005 by IxIoehr(2005) in AISI.	SN 2005cs was discovered on 29th June 2005 by \citet{discov05cs} in M51.
Due to its very carly cliscovery and proximity. this object is very well-sampled. and. studied in detail.	Due to its very early discovery and proximity, this object is very well-sampled and studied in detail.
is a uncerluminous. low-energy. Ni-poor SN that had a low-mass progenitor (see Table 1. for references).	It is a underluminous, low-energy, Ni-poor SN that had a low-mass progenitor (see Table \ref{physdata} for references).
14 spectra of Pastorelloetal.(2006). and. Pastorelloctal. were used. which were obtained. between days. [3 and\|61.	14 spectra of \citet{pastorello05csI} and \citet{pastorello05csII} were used, which were obtained between days +3 and+61.
SN 2006bp was discovered on 9th April 2006 by (Nakano in NCC 3953. also in àvery early phase.	SN 2006bp was discovered on 9th April 2006 by \citep{discov06bp} in NGC 3953, also in avery early phase.
We used 11 spectra of Quimbyctal. (2007)... downloaded from SUSPECT. covering the period between \| 5 ancl \| 72 days.	We used 11 spectra of \citet{quimby06bp}, , downloaded from , covering the period between $+$ 5 and $+$ 72 days.
Le. with PBL (<100m in our study) - for our evaluation.	i.e. with PBL $<100$ m in our study) - for our evaluation.
The free atmosphere forecast will be compared with the observations [rom Multi-Aperture Scintillation Sensor (MASS) which is desigued to measure the seeiug iu (ree atinospliere (Tokovinin2002).	The free atmosphere forecast will be compared with the observations from Multi-Aperture Scintillation Sensor (MASS) which is designed to measure the seeing in free atmosphere \citep{tok02}.
. We collect cloud cover and seeing observations for a total of nine sites [roi Central/East Asia. Hawaliaxd Central/South America iithe period of January 2008 to December 2009.	We collect cloud cover and seeing observations for a total of nine sites from Central/East Asia, Hawaii and Central/South America in the period of January 2008 to December 2009.
The respective information of each site includiug tle PBL top limit and P(A) used in the GFS/AXP model are isted i Table 3..	The respective information of each site including the PBL top limit and $P(h)$ used in the GFS/AXP model are listed in Table \ref{tbl-3}.
As these observations are a] inace in sequence with sampling frequency around 1-1.510in except anmsha1 and Lulin (whicl will be «eal separately aud will be described below). tley are firstly jyocessed to match the time iilerva of the GES model output (which is 3h).	As these observations are all made in sequence with sampling frequency around 1-1.5min except Nanshan and Lulin (which will be deal separately and will be described below), they are firstly processed to match the time interval of the GFS model output (which is 3h).
To οιsure that tLe observation is representative ii the corresponding interval. we set a 1ninunuim claa poluts of 100 or each ierval.	To ensure that the observation is representative in the corresponding interval, we set a minimum data points of 100 for each interval.
A sampli1οD> requency of 1-1σα correspolrds 120-180 data poins per 3h. so the nininiuim hreshold of 100 Isi 'easolable.	A sampling frequency of 1-1.5min corresponds 120-180 data points per 3h, so the minimum threshold of 100 is reasonable.
Asal seeiug observatious are obtained either by Diflerentia Image Motion NdLitor (DIMM) or MASS. hey cau be used without further reduction since they eive the rueasuremenut of eg cirectly.	As all seeing observations are obtained either by Differential Image Motion Monitor (DIMM) or MASS, they can be used without further reduction since they give the measurement of $\epsilon_{0}$ directly.
On the ojer haud. the clouc cover observations from Paralal. Naushanu aud Lulin are obtained with different tools. so they iust be reduced to the same definilion with the GFS moclel output before couparing them with the later.	On the other hand, the cloud cover observations from Paranal, Nanshan and Lulin are obtained with different tools, so they must be reduced to the same definition with the GFS model output before comparing them with the later.
The reduction procedwe is descri)ed below.	The reduction procedure is described below.
The cloud sensor instaled at Paranal determines the sky coiditiou by measuring the flux variation of a star.	The cloud sensor installed at Paranal determines the sky condition by measuring the flux variation of a star.
The seusr graph will suggests a possible "clouly” condition when the root-mean-square (RMS) of the fhx variation is larger than	The sensor graph will suggests a possible “cloudy” condition when the root-mean-square (RMS) of the flux variation is larger than.
For Naushau. the ope‘ation log is used to verily the sky condition.	For Nanshan, the operation log is used to verify the sky condition.
The operation log includes the image sequence log ane Lobservers notes.	The operation log includes the image sequence log and observer's notes.
We divide each night into two “paris”: evening and morning.	We divide each night into two “parts”: evening and morning.
A “part” with rouglly >90% observable time (with images taken aud indication of eood observing condition [rom the observer) would be marked as “clear”. otherwise it would be marked as "cloudy.	A “part” with roughly $>50$ observable time (with images taken and indication of good observing condition from the observer) would be marked as “clear”, otherwise it would be marked as “cloudy”.
edee of our resolved RGB stars 1s wld. if 00.004L. it is around 1.2. aud. if 00.0001. it is around 1.0.	edge of our resolved RGB stars is $\,\simeq\,$ 1.4, if 0.004, it is around 1.2, and, if 0.0004, it is around 1.0.
If we consider as RGB stars all those with c 2L25 and redder than these blue edges. «'e find that 3511. 2398 or 1227 Clor 00.008. 0.00L or 0.000L. respectively) over 0221 stars are older tiui lL Gyr.	If we consider as RGB stars all those with $\,\geq\,$ 24.5 and redder than these blue edges, we find that 3511, 2398 or 1227 (for 0.008, 0.004, or 0.0004, respectively) over 9224 stars are older than 1 Gyr.
In other words. thauks to the higher resolving power of the HST. we have been able to measure also iuη the innerη x ppc2 the oldest population. which was too [aint to be resolvabe with grouud-based observations.	In other words, thanks to the higher resolving power of the HST, we have been able to measure also in the inner $\,\times\,$ $^2$ the oldest population, which was too faint to be resolvable with ground-based observations.
The diagramse of Fie.	The diagrams of Fig.
e I7. confirm the existence of an age gradient.	\ref{cm_vi} confirm the existence of an age gradient.
We find that: As mentioned i1 the Introduction. previous stuclies already pointed out that NGC 1705 has a coniposite populaion (Quillen et al.	We find that: As mentioned in the Introduction, previous studies already pointed out that NGC 1705 has a composite population (Quillen et al.
1995). with an age gradient. (MEDC) lise the one here described.	1995), with an age gradient (MFDC) like the one here described.
This is. however. the first evidence from clirect analysis of the resolved stela' populatious.	This is, however, the first evidence from direct analysis of the resolved stellar populations.
The same kind of spatial segregatioue of the vouugerDm stars has been found also in [9]her late-type dwarls (e.g.. NGC 1569. G98: I Zw 18. ATC: VIL Zw 103. Sehulte-Lacdbeck et al.	The same kind of spatial segregation of the younger stars has been found also in other late-type dwarfs (e.g., NGC 1569, G98; I Zw 18, ATG; VII Zw 403, Schulte-Ladbeck et al.
19092).	1999a).
Iu stummary. NGC 1705 is definitely not a st:uwburst galaxy. since it cloes uot appear to have had auy conspleuous SF activity in the last [ew AMyrs. neither in the fiekl nor iu he SSC.	In summary, NGC 1705 is definitely not a starburst galaxy, since it does not appear to have had any conspicuous SF activity in the last few Myrs, neither in the field nor in the SSC.
Our CMDs show that the last episode of significant SF in the fiekl has occurred around 1-20 Myr ago. Isreceded by several οἱier episodes or by a continuous activity.	Our CMDs show that the last episode of significant SF in the field has occurred around 10–20 Myr ago, preceded by several other episodes or by a continuous activity.
To uuderstaid whetler or not aly “these episodes can ye considered as a real burst (Le. short aud interse SF activily). ancl how long the quiescent phases (if any) have lasted. oue must perform a more sopUsticatec analysis.	To understand whether or not any of these episodes can be considered as a real burst (i.e. short and intense SF activity), and how long the quiescent phases (if any) have lasted, one must perform a more sophisticated analysis.
We 'e working oi this atalysis using the method of synthetic CMDs described by Tosi et al. (	We are working on this analysis using the method of synthetic CMDs described by Tosi et al. (
1991) xd COS. and the results will be presented in a ortheoming paper (Anuibali et al.	1991) and G98, and the results will be presented in a forthcoming paper (Annibali et al.
in preparatiou).	in preparation).
The impressive morphology of the H4 images of the galaxy. counated by loops and ares ipparently centered on the brightest SSC. ane ie correspondiug kinematics have been studied w MEDC.	The impressive morphology of the $_{\alpha}$ images of the galaxy, dominated by loops and arcs apparently centered on the brightest SSC, and the corresponding kinematics have been studied by MFDC.
They interoreted the loops and ares as hot expaudiug bubbles energized by supernova jecta and stellar winds from its uucleus. with esimatecd expansion timescale of 9 Myr.	They interpreted the loops and arcs as hot expanding bubbles energized by supernova ejecta and stellar winds from its nucleus, with estimated expansion timescale of $\sim$ 9 Myr.
Recent USE observations of NCC 1705 (Heckman οἱ al.	Recent FUSE observations of NGC 1705 (Heckman et al.
2001) also show that the superbubbles will most ire!‘obably blow out of the galaxy. aid remove a sieuilicaut fraction of its metals.	2001) also show that the superbubbles will most probably blow out of the galaxy and remove a significant fraction of its metals.
Sahu Blades 1997) agreed with this scenario of conspleuous [n]oOalactic wiud triggered by supernova explosions. ithe basis of HST UV spectra.	Sahu Blades (1997) agreed with this scenario of conspicuous galactic wind triggered by supernova explosions, on the basis of HST UV spectra.
Their data confirmed the 510 kin \|. velocity of the supershell	Their data confirmed the 540 km $^{-1}$ velocity of the supershell
For galaxies for which a J baud nuage is also available. we display in Figs.	For galaxies for which a J band image is also available, we display in Figs.
Ed-28d the ratio of the J to I images (after background. subtraction). median filtered by 345 pixels in order to enhance the S/N of the outer regions.	1d-28d the ratio of the J to K' images (after background subtraction), median filtered by $\times$ 3 pixels in order to enhance the S/N of the outer regions.
Color images will allow a more complete description of the previously detected morphological features iu terms of the contribution of dust and/or star forming events.	Color images will allow a more complete description of the previously detected morphological features in terms of the contribution of dust and/or star forming events.
We uote that in the cases where the background is not flat. somewhat structured color unages are obtained (galaxies with nou-Hat remaining backgrounds are imdicated with an asterisk in col 1 of Table 2)).	We note that in the cases where the background is not flat, somewhat structured color images are obtained (galaxies with non-flat remaining backgrounds are indicated with an asterisk in column 1 of Table \ref{journal}) ).
AMaeuitudes iuteerated im cieuluapertures and J- color eradicuts were obtained frou the curves of erowthn as in Mürrquez Moles (1996).	Magnitudes integrated in circularapertures and J-K' color gradients were obtained from the curves of growth as in Márrquez Moles (1996).
The apertures aud the corresponding maenitudes are eivon in Table 3..	The apertures and the corresponding magnitudes are given in Table \ref{decomposition}.
Color eracdicuts are shown in Figs.	Color gradients are shown in Figs.
Lh-2s8h for galaxies for which we have data in the two bands.	1h-28h for galaxies for which we have data in the two bands.
The backgrouud subtraction was done as described above aud the same caveat applies. in the seuse that some color eradieut could be artificially produced for the galaxies with mon-flat backerounds.	The background subtraction was done as described above and the same caveat applies, in the sense that some color gradient could be artificially produced for the galaxies with non-flat backgrounds.
Iu Table [\| we eive the magnitudes that we measure for 9 ealaxies. together with J. I and/or I& maenitudes found in the literature.	In Table \ref{compara} we give the magnitudes that we measure for 9 galaxies, together with J, K' and/or K magnitudes found in the literature.
A straightforward comparison between our J values aud those reported iu previous works (for the same apertures) gives Jas ρε=O12£0.25.	A straightforward comparison between our J values and those reported in previous works (for the same apertures) gives: $_{us} -$ $_{other} = -0.42 \pm 0.25$.
Oulv 1 galaxies in our sample have been observed in Ij by Mulchaeyv. et al. (	Only 4 galaxies in our sample have been observed in K' by Mulchaey et al. (
1997): we fud: Ius \nuchace0.015+ 0.63.	1997); we find: $_{us} -$ $_{\rm Mulchaey} = 0.015 \pm 0.63$ .
If we exclude NGC 6890 (see section 5.17) we obtain: KR, Iba0.304.0.36.	If we exclude NGC 6890 (see section 5.17) we obtain: $_{us} -$ $_{\rm Mulchaey} = -0.30 \pm 0.36$.
For 10 ealaxies aud a total of 12 measurements. we find: IN, lone=o0.12£0.32.	For 10 galaxies and a total of 12 measurements, we find: $_{us} -$ $_{other} = -0.42 \pm 0.32$.
This value is comparable to the difference between ypiegaec and IS; 0.55d: 0.19).	This value is comparable to the difference between $_{\rm Mulchaey}$ and $_{other}$ $-0.55 \pm 0.49$ ).
These differences in the zero points are not critical for our purposes aud will uot be discussed further.	These differences in the zero points are not critical for our purposes and will not be discussed further.
Ta any case. we notice that the differences in the colour iudexes J-IKC are withiu the errors.	In any case, we notice that the differences in the colour indexes J-K' are within the errors.
Iu Table 5 we eive the parameters we determine for primary aud secondary bars in the sample galaxies;	In Table \ref{bars} we give the parameters we determine for primary and secondary bars in the sample galaxies.
We specity what of the methods above have been used to detect the presence of secondary bars (or primary bars when detected for the first time). and if evidences are found iu available UST images.	We specify what of the methods above have been used to detect the presence of secondary bars (or primary bars when detected for the first time), and if evidences are found in available HST images.
We note that. excepting the secondary bar of IC. 2510. only secu in the sharp-divided nuage (just mareinally in the PA-e ot). the secondary bars are obtained iu at least two methods.	We note that, excepting the secondary bar of IC 2510, only seen in the sharp-divided image (just marginally in the $\epsilon$ plot), the secondary bars are obtained in at least two methods.
The I. nuage in Fig.	The K' image in Fig.

. The drawback of these approaches is that the ruilcing of tailored. model atmospheres can be very. time-consuming and needs much more computing power.

The drawback of these approaches is that the building of tailored model atmospheres can be very time-consuming and needs much more computing power.

Also. in order to compute reliable models. the input observed spectra need to have sulliciently high S/N and spectral resolution.

Also, in order to compute reliable models, the input observed spectra need to have sufficiently high S/N and spectral resolution.

The more recently developed: Standardized Canelle AMethocl (SCAL) (Llamuy&Pinto2002:Lamuyοἱal.2003) relies on the empirical correlationbetween the measured expansion velocity and the luminosity in optical (Nugentetal.2006:Poznanski2009). or in NI. (Maguireetal.2010a)— bands in the micelle of plateau phase. αἲ [50 days after explosion.

The more recently developed Standardized Candle Method (SCM) \citep{hamuy_pinto_scm, hamuy_scm} relies on the empirical correlationbetween the measured expansion velocity and the luminosity in optical \citep{nugent06, poznanski09} or in NIR \citep{maguirescm} bands in the middle of plateau phase, at +50 days after explosion.

“Phis method. requires. less extensive input. data. but needs a larger sample of Tvpe LL SNe with independently known distances in order to calibrate the empirical correlation.

This method requires less extensive input data, but needs a larger sample of Type II SNe with independently known distances in order to calibrate the empirical correlation.

Because this is basically a wxhotometrice method which compares apparent and absolute magnitudes. SCAL is more sensitive to interstellar reddening han EDPM/SIEXM. as mentioned above.

Because this is basically a photometric method which compares apparent and absolute magnitudes, SCM is more sensitive to interstellar reddening than EPM/SEAM, as mentioned above.

Beside photometry. both IEPMSIZAM and SCM need information on the expansion velocity at the photosphere (Pye).

Beside photometry, both EPM/SEAM and SCM need information on the expansion velocity at the photosphere $v_{phot}$ ).

At first. EPM seems to be more challenging. because it requires multi-epoch observations. while SCAL needs only he velocity at a certain epoch. at. |50 days after explosion (eso).

At first, EPM seems to be more challenging, because it requires multi-epoch observations, while SCM needs only the velocity at a certain epoch, at +50 days after explosion $v_{50}$ ).

However. direct measurement of eso is possible only in he case of precise timing of the spectroscopic observation. which is rarely achievable.

However, direct measurement of $v_{50}$ is possible only in the case of precise timing of the spectroscopic observation, which is rarely achievable.

The correct. determination of Pii; is not trivial.

The correct determination of $v_{phot}$ is not trivial.

As SNe expand homologously (er. where e is the velocity of a given laver and r is its distance from the center). in most cases it ds dillicult to derive a unique velocity [rom the observed. spectral features.

As SNe expand homologously $v \sim r$, where $v$ is the velocity of a given layer and $r$ is its distance from the center), in most cases it is difficult to derive a unique velocity from the observed spectral features.

The most. [frequently Followed approach relies on measuring the Doppler-shift of the absorption minimum of certain spectral features (mostly A5169 or 11:2).

The most frequently followed approach relies on measuring the Doppler-shift of the absorption minimum of certain spectral features (mostly $\lambda 5169$ or $\beta$ ).

Another possibility is the computation of the eross-correlation between the SN spectrum and a set of spectral templates. consisting of either observed or model spectra. with known velocities.

Another possibility is the computation of the cross-correlation between the SN spectrum and a set of spectral templates, consisting of either observed or model spectra, with known velocities.

Phe third method is building a [Full NUTE model (like PILOENIN or CAIPCIEN) for à given SN spectrum. and. adopting the theoretical ρω from the rest-Litting moclel.

The third method is building a full NLTE model (like PHOENIX or CMFGEN) for a given SN spectrum, and adopting the theoretical $v_{phot}$ from the best-fitting model.

The aim of this paper is to present a similar. but ess computation-demanding approach to assign velocities ο observed SN spectra.

The aim of this paper is to present a similar, but less computation-demanding approach to assign velocities to observed SN spectra.

We apply the simple parametrized code (Fisher1999:Llatanoctal.1999). to model the observed spectra with an approximative. but self-consistent reatment of the formation of lines in the extended. 10mologously. expanding SN atmosphere.

We apply the simple parametrized code \citep{fisher99, hatano99} to model the observed spectra with an approximative, but self-consistent treatment of the formation of lines in the extended, homologously expanding SN atmosphere.

To illustrate the applicability ofSYNOW. we construct ancl fit. parametrized models to a sample of five wellcobserved. nearby Type H-P SNe that have a series of high S/N spectra publicly available.

To illustrate the applicability of, we construct and fit parametrized models to a sample of five well-observed, nearby Type II-P SNe that have a series of high S/N spectra publicly available.

We also check ancl re-calibrate some empirical correlations between velocities derived. from various methods.

We also check and re-calibrate some empirical correlations between velocities derived from various methods.

The deseription of the observational sample is given in δεο,

The description of the observational sample is given in \ref{sec_data}.

In §?? we first review the cdillerent. velocity measurement techniques for SNe U-P. then we present the details of the application of SYNOW models (§??7)).

In \ref{sec_issue} we first review the different velocity measurement techniques for SNe II-P, then we present the details of the application of models \ref{sec_synow}) ).

Phe results are collected ancl discussed in $?? and 8?7?.. while the implications of the results for the clistance measurements are in $??..

The results are collected and discussed in \ref{sec_results} and \ref{sec_disc}, while the implications of the results for the distance measurements are in \ref{sec_dist}.

We draw our conclusions in 8?7..

We draw our conclusions in \ref{sec_concl}.

Our sample contains S1 plateau phasespectra of five objects. SNe 1999em. 2004dj. 2004et. 2005es and 2006bp. respectively.

Our sample contains 81 plateau phase–spectra of five objects, SNe 1999em, 2004dj, 2004et, 2005cs and 2006bp, respectively.

All five SNe are wellobserved objects. they have been stucied in detail. and show a wide variety in their physical properties (see Table 1)).

All five SNe are well-observed objects, they have been studied in detail, and show a wide variety in their physical properties (see Table \ref{physdata}) ).

SN 1999em was discovered. on 29th October 1999. by Li(19900) in NGC 1637 at a very carly phase.

SN 1999em was discovered on 29th October 1999 by \citet{discov99em} in NGC 1637 at a very early phase.

Ht is a very. well-observed. well-studied object The explosion date was determined as 2451477.0€2 JD (Leonardοἱal.2002:Παινetal.2001:Dessart&LHillier 2006)...

It is a very well-observed, well-studied object The explosion date was determined as $2451477.0 \pm 2$ JD \citep{leonard99em, hamuy99em, dessart2006}. .

We used the spectra of Leonardetal.(2002). and. Hamus.ctal.(2001) downloaded. from the database.

We used the spectra of \citet{leonard99em} and \citet{hamuy99em} downloaded from the database.

Our. saniple contains 22 spectra. covering the first SO days of the plateu phase.

Our sample contains 22 spectra, covering the first 80 days of the plateu phase.

SN 2004dj was discovered on 31st July 2004 by Iagaki (Nakanoetal.2004) in à voung. massive cluster Sancdage-96 of NGC 2403. about 1: month after explosion.

SN 2004dj was discovered on 31st July 2004 by Itagaki \citep{discov04dj} in a young, massive cluster Sandage-96 of NGC 2403, about 1 month after explosion.

We used the spectra taken by Vinkóetal.(2006)..

We used the spectra taken by \citet{vinko04dj}.

Due to the lack of observed spectrophotometric standards. the lux-calibration of those spectra was inferior. but this is not a major concern when only velocities are to be determined.

Due to the lack of observed spectrophotometric standards, the flux-calibration of those spectra was inferior, but this is not a major concern when only velocities are to be determined.

We included. 12 spectra taken between |47 ancl |100 days after explosion in our sample.

We included 12 spectra taken between +47 and +100 days after explosion in our sample.

SN 2004et was discovered by Moretti (Zwitter.Murani&Moretti2004) on 27th September 2004 in NGC 6946.

SN 2004et was discovered by Moretti \citep*{discov04et} on 27th September 2004 in NGC 6946.

The spectra of Sahuetal.(2006). (downloaded. (rom SUSPECT) and Maguireetal.(2010b) were used. together. with 6 previously unpublished earlv-phase spectra taken with the l.8S-m telescope at DDO (see the Appendix ??7)).

The spectra of \citet{sahu04et} (downloaded from ) and \citet{maguire04et} were used, together with 6 previously unpublished early-phase spectra taken with the 1.88-m telescope at DDO (see the Appendix \ref{appendix_a}) ).

The 22 spectra cover the period of | 11 to | 104 days after explosion.

The 22 spectra cover the period of $+$ 11 to $+$ 104 days after explosion.

SN 2005cs was discovered on 29th June 2005 by IxIoehr(2005) in AISI.

SN 2005cs was discovered on 29th June 2005 by \citet{discov05cs} in M51.

Due to its very carly cliscovery and proximity. this object is very well-sampled. and. studied in detail.

Due to its very early discovery and proximity, this object is very well-sampled and studied in detail.

is a uncerluminous. low-energy. Ni-poor SN that had a low-mass progenitor (see Table 1. for references).

It is a underluminous, low-energy, Ni-poor SN that had a low-mass progenitor (see Table \ref{physdata} for references).

14 spectra of Pastorelloetal.(2006). and. Pastorelloctal. were used. which were obtained. between days. [3 and|61.

14 spectra of \citet{pastorello05csI} and \citet{pastorello05csII} were used, which were obtained between days +3 and+61.

SN 2006bp was discovered on 9th April 2006 by (Nakano in NCC 3953. also in àvery early phase.

SN 2006bp was discovered on 9th April 2006 by \citep{discov06bp} in NGC 3953, also in avery early phase.

We used 11 spectra of Quimbyctal. (2007)... downloaded from SUSPECT. covering the period between | 5 ancl | 72 days.

We used 11 spectra of \citet{quimby06bp}, , downloaded from , covering the period between $+$ 5 and $+$ 72 days.

Le. with PBL (<100m in our study) - for our evaluation.

i.e. with PBL $<100$ m in our study) - for our evaluation.

The free atmosphere forecast will be compared with the observations [rom Multi-Aperture Scintillation Sensor (MASS) which is desigued to measure the seeiug iu (ree atinospliere (Tokovinin2002).

The free atmosphere forecast will be compared with the observations from Multi-Aperture Scintillation Sensor (MASS) which is designed to measure the seeing in free atmosphere \citep{tok02}.

. We collect cloud cover and seeing observations for a total of nine sites [roi Central/East Asia. Hawaliaxd Central/South America iithe period of January 2008 to December 2009.

We collect cloud cover and seeing observations for a total of nine sites from Central/East Asia, Hawaii and Central/South America in the period of January 2008 to December 2009.

The respective information of each site includiug tle PBL top limit and P(A) used in the GFS/AXP model are isted i Table 3..

The respective information of each site including the PBL top limit and $P(h)$ used in the GFS/AXP model are listed in Table \ref{tbl-3}.

As these observations are a] inace in sequence with sampling frequency around 1-1.510in except anmsha1 and Lulin (whicl will be «eal separately aud will be described below). tley are firstly jyocessed to match the time iilerva of the GES model output (which is 3h).

As these observations are all made in sequence with sampling frequency around 1-1.5min except Nanshan and Lulin (which will be deal separately and will be described below), they are firstly processed to match the time interval of the GFS model output (which is 3h).

To οιsure that tLe observation is representative ii the corresponding interval. we set a 1ninunuim claa poluts of 100 or each ierval.

To ensure that the observation is representative in the corresponding interval, we set a minimum data points of 100 for each interval.

A sampli1οD> requency of 1-1σα correspolrds 120-180 data poins per 3h. so the nininiuim hreshold of 100 Isi 'easolable.

A sampling frequency of 1-1.5min corresponds 120-180 data points per 3h, so the minimum threshold of 100 is reasonable.

Asal seeiug observatious are obtained either by Diflerentia Image Motion NdLitor (DIMM) or MASS. hey cau be used without further reduction since they eive the rueasuremenut of eg cirectly.

As all seeing observations are obtained either by Differential Image Motion Monitor (DIMM) or MASS, they can be used without further reduction since they give the measurement of $\epsilon_{0}$ directly.

On the ojer haud. the clouc cover observations from Paralal. Naushanu aud Lulin are obtained with different tools. so they iust be reduced to the same definilion with the GFS moclel output before couparing them with the later.

On the other hand, the cloud cover observations from Paranal, Nanshan and Lulin are obtained with different tools, so they must be reduced to the same definition with the GFS model output before comparing them with the later.

The reduction procedwe is descri)ed below.

The reduction procedure is described below.

The cloud sensor instaled at Paranal determines the sky coiditiou by measuring the flux variation of a star.

The cloud sensor installed at Paranal determines the sky condition by measuring the flux variation of a star.

The seusr graph will suggests a possible "clouly” condition when the root-mean-square (RMS) of the fhx variation is larger than

The sensor graph will suggests a possible “cloudy” condition when the root-mean-square (RMS) of the flux variation is larger than.

For Naushau. the ope‘ation log is used to verily the sky condition.

For Nanshan, the operation log is used to verify the sky condition.

The operation log includes the image sequence log ane Lobservers notes.

The operation log includes the image sequence log and observer's notes.

We divide each night into two “paris”: evening and morning.

We divide each night into two “parts”: evening and morning.

A “part” with rouglly >90% observable time (with images taken aud indication of eood observing condition [rom the observer) would be marked as “clear”. otherwise it would be marked as "cloudy.

A “part” with roughly $>50$ observable time (with images taken and indication of good observing condition from the observer) would be marked as “clear”, otherwise it would be marked as “cloudy”.

edee of our resolved RGB stars 1s wld. if 00.004L. it is around 1.2. aud. if 00.0001. it is around 1.0.

edge of our resolved RGB stars is $\,\simeq\,$ 1.4, if 0.004, it is around 1.2, and, if 0.0004, it is around 1.0.

If we consider as RGB stars all those with c 2L25 and redder than these blue edges. «'e find that 3511. 2398 or 1227 Clor 00.008. 0.00L or 0.000L. respectively) over 0221 stars are older tiui lL Gyr.

If we consider as RGB stars all those with $\,\geq\,$ 24.5 and redder than these blue edges, we find that 3511, 2398 or 1227 (for 0.008, 0.004, or 0.0004, respectively) over 9224 stars are older than 1 Gyr.

In other words. thauks to the higher resolving power of the HST. we have been able to measure also iuη the innerη x ppc2 the oldest population. which was too [aint to be resolvabe with grouud-based observations.

In other words, thanks to the higher resolving power of the HST, we have been able to measure also in the inner $\,\times\,$ $^2$ the oldest population, which was too faint to be resolvable with ground-based observations.

The diagramse of Fie.

The diagrams of Fig.

e I7. confirm the existence of an age gradient.

\ref{cm_vi} confirm the existence of an age gradient.

We find that: As mentioned i1 the Introduction. previous stuclies already pointed out that NGC 1705 has a coniposite populaion (Quillen et al.

We find that: As mentioned in the Introduction, previous studies already pointed out that NGC 1705 has a composite population (Quillen et al.

1995). with an age gradient. (MEDC) lise the one here described.

1995), with an age gradient (MFDC) like the one here described.

This is. however. the first evidence from clirect analysis of the resolved stela' populatious.

This is, however, the first evidence from direct analysis of the resolved stellar populations.

The same kind of spatial segregatioue of the vouugerDm stars has been found also in [9]her late-type dwarls (e.g.. NGC 1569. G98: I Zw 18. ATC: VIL Zw 103. Sehulte-Lacdbeck et al.

The same kind of spatial segregation of the younger stars has been found also in other late-type dwarfs (e.g., NGC 1569, G98; I Zw 18, ATG; VII Zw 403, Schulte-Ladbeck et al.

19092).

1999a).

Iu stummary. NGC 1705 is definitely not a st:uwburst galaxy. since it cloes uot appear to have had auy conspleuous SF activity in the last [ew AMyrs. neither in the fiekl nor iu he SSC.

In summary, NGC 1705 is definitely not a starburst galaxy, since it does not appear to have had any conspicuous SF activity in the last few Myrs, neither in the field nor in the SSC.

Our CMDs show that the last episode of significant SF in the fiekl has occurred around 1-20 Myr ago. Isreceded by several οἱier episodes or by a continuous activity.

Our CMDs show that the last episode of significant SF in the field has occurred around 10–20 Myr ago, preceded by several other episodes or by a continuous activity.

To uuderstaid whetler or not aly “these episodes can ye considered as a real burst (Le. short aud interse SF activily). ancl how long the quiescent phases (if any) have lasted. oue must perform a more sopUsticatec analysis.

To understand whether or not any of these episodes can be considered as a real burst (i.e. short and intense SF activity), and how long the quiescent phases (if any) have lasted, one must perform a more sophisticated analysis.

We 'e working oi this atalysis using the method of synthetic CMDs described by Tosi et al. (

We are working on this analysis using the method of synthetic CMDs described by Tosi et al. (

1991) xd COS. and the results will be presented in a ortheoming paper (Anuibali et al.

1991) and G98, and the results will be presented in a forthcoming paper (Annibali et al.

in preparatiou).

in preparation).

The impressive morphology of the H4 images of the galaxy. counated by loops and ares ipparently centered on the brightest SSC. ane ie correspondiug kinematics have been studied w MEDC.

The impressive morphology of the $_{\alpha}$ images of the galaxy, dominated by loops and arcs apparently centered on the brightest SSC, and the corresponding kinematics have been studied by MFDC.

They interoreted the loops and ares as hot expaudiug bubbles energized by supernova jecta and stellar winds from its uucleus. with esimatecd expansion timescale of 9 Myr.

They interpreted the loops and arcs as hot expanding bubbles energized by supernova ejecta and stellar winds from its nucleus, with estimated expansion timescale of $\sim$ 9 Myr.

Recent USE observations of NCC 1705 (Heckman οἱ al.

Recent FUSE observations of NGC 1705 (Heckman et al.

2001) also show that the superbubbles will most ire!‘obably blow out of the galaxy. aid remove a sieuilicaut fraction of its metals.

2001) also show that the superbubbles will most probably blow out of the galaxy and remove a significant fraction of its metals.

Sahu Blades 1997) agreed with this scenario of conspleuous [n]oOalactic wiud triggered by supernova explosions. ithe basis of HST UV spectra.

Sahu Blades (1997) agreed with this scenario of conspicuous galactic wind triggered by supernova explosions, on the basis of HST UV spectra.

Their data confirmed the 510 kin |. velocity of the supershell

Their data confirmed the 540 km $^{-1}$ velocity of the supershell

For galaxies for which a J baud nuage is also available. we display in Figs.

For galaxies for which a J band image is also available, we display in Figs.

Ed-28d the ratio of the J to I images (after background. subtraction). median filtered by 345 pixels in order to enhance the S/N of the outer regions.

1d-28d the ratio of the J to K' images (after background subtraction), median filtered by $\times$ 3 pixels in order to enhance the S/N of the outer regions.

Color images will allow a more complete description of the previously detected morphological features iu terms of the contribution of dust and/or star forming events.

Color images will allow a more complete description of the previously detected morphological features in terms of the contribution of dust and/or star forming events.

We uote that in the cases where the background is not flat. somewhat structured color unages are obtained (galaxies with nou-Hat remaining backgrounds are imdicated with an asterisk in col 1 of Table 2)).

We note that in the cases where the background is not flat, somewhat structured color images are obtained (galaxies with non-flat remaining backgrounds are indicated with an asterisk in column 1 of Table \ref{journal}) ).

AMaeuitudes iuteerated im cieuluapertures and J- color eradicuts were obtained frou the curves of erowthn as in Mürrquez Moles (1996).

Magnitudes integrated in circularapertures and J-K' color gradients were obtained from the curves of growth as in Márrquez Moles (1996).

The apertures aud the corresponding maenitudes are eivon in Table 3..

The apertures and the corresponding magnitudes are given in Table \ref{decomposition}.

Color eracdicuts are shown in Figs.

Color gradients are shown in Figs.

Lh-2s8h for galaxies for which we have data in the two bands.

1h-28h for galaxies for which we have data in the two bands.

The backgrouud subtraction was done as described above aud the same caveat applies. in the seuse that some color eradieut could be artificially produced for the galaxies with mon-flat backerounds.

The background subtraction was done as described above and the same caveat applies, in the sense that some color gradient could be artificially produced for the galaxies with non-flat backgrounds.

Iu Table [| we eive the magnitudes that we measure for 9 ealaxies. together with J. I and/or I& maenitudes found in the literature.

In Table \ref{compara} we give the magnitudes that we measure for 9 galaxies, together with J, K' and/or K magnitudes found in the literature.

A straightforward comparison between our J values aud those reported iu previous works (for the same apertures) gives Jas ρε=O12£0.25.

A straightforward comparison between our J values and those reported in previous works (for the same apertures) gives: $_{us} -$ $_{other} = -0.42 \pm 0.25$.

Oulv 1 galaxies in our sample have been observed in Ij by Mulchaeyv. et al. (

Only 4 galaxies in our sample have been observed in K' by Mulchaey et al. (

1997): we fud: Ius \nuchace0.015+ 0.63.

1997); we find: $_{us} -$ $_{\rm Mulchaey} = 0.015 \pm 0.63$ .

If we exclude NGC 6890 (see section 5.17) we obtain: KR, Iba0.304.0.36.

If we exclude NGC 6890 (see section 5.17) we obtain: $_{us} -$ $_{\rm Mulchaey} = -0.30 \pm 0.36$.

For 10 ealaxies aud a total of 12 measurements. we find: IN, lone=o0.12£0.32.

For 10 galaxies and a total of 12 measurements, we find: $_{us} -$ $_{other} = -0.42 \pm 0.32$.

This value is comparable to the difference between ypiegaec and IS; 0.55d: 0.19).

This value is comparable to the difference between $_{\rm Mulchaey}$ and $_{other}$ $-0.55 \pm 0.49$ ).

These differences in the zero points are not critical for our purposes aud will uot be discussed further.

These differences in the zero points are not critical for our purposes and will not be discussed further.

Ta any case. we notice that the differences in the colour iudexes J-IKC are withiu the errors.

In any case, we notice that the differences in the colour indexes J-K' are within the errors.

Iu Table 5 we eive the parameters we determine for primary aud secondary bars in the sample galaxies;

In Table \ref{bars} we give the parameters we determine for primary and secondary bars in the sample galaxies.

We specity what of the methods above have been used to detect the presence of secondary bars (or primary bars when detected for the first time). and if evidences are found iu available UST images.

We specify what of the methods above have been used to detect the presence of secondary bars (or primary bars when detected for the first time), and if evidences are found in available HST images.

We note that. excepting the secondary bar of IC. 2510. only secu in the sharp-divided nuage (just mareinally in the PA-e ot). the secondary bars are obtained iu at least two methods.

We note that, excepting the secondary bar of IC 2510, only seen in the sharp-divided image (just marginally in the $\epsilon$ plot), the secondary bars are obtained in at least two methods.

The I. nuage in Fig.

The K' image in Fig.