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Mediuu-resolutiou. optical spectra of the objects im the XRT eror circle (sce Fig. | Medium-resolution optical spectra of the objects in the /XRT error circle (see Fig. |
1) were acquired between 0:21 aud 06:23 UT of 21 June 2007. and between03:3 and 0101 UT of 21 June 2007 with the 3.6-uctre ESO telescope located in La Silla (Chile). | 1) were acquired between 04:24 and 06:23 UT of 21 June 2007, and between03:34 and 04:04 UT of 24 June 2007 with the 3.6-metre ESO telescope located in La Silla (Chile). |
This telescope carrice the EFOSC?2 instrmucut. equipped with a «2018 pixel Loral/Lesser CCD. | This telescope carried the EFOSC2 instrument, equipped with a $\times$ 2048 pixel Loral/Lesser CCD. |
The use of erating #1153 anc a slit of 170 provided a 36859315 nonminual spectra coverage, | The use of grating 13 and a slit of $\farcs$ 0 provided a 3685–9315 nominal spectral coverage. |
This setup gave a dispersion of 2.8 A//pix. | This setup gave a dispersion of 2.8 /pix. |
The spectra (Fig. | The spectra (Fig. |
2). after correction for cosmic-ray rejection. bias and flat-field. were optimally extractec | 2), after correction for cosmic-ray rejection, bias and flat-field, were optimally extracted |
Including these realistic errors one would expect about 52.5 percent of the stars to have proper motions within fortyv-five degrees of the radial direction. so about 1.600. stars would be required to detect this asymmetry at the two-siema level if the astrometric errors in current instruments are included. | Including these realistic errors one would expect about 52.5 percent of the stars to have proper motions within forty-five degrees of the radial direction, so about 1,600 stars would be required to detect this asymmetry at the two-sigma level if the astrometric errors in current instruments are included. |
I white cbwarfs do indeed receive. kicks as they form. both their racial and velocity distributions will appear cillerent from their progenitors and the older white chvarls whose distribution has relaxed. | If white dwarfs do indeed receive kicks as they form, both their radial and velocity distributions will appear different from their progenitors and the older white dwarfs whose distribution has relaxed. |
The easier of these ellects to measure is the change in the racial distribution of the stars (?): however. there are two other lines of evidence that might be present in astrometric measurements of globular clusters. | The easier of these effects to measure is the change in the radial distribution of the stars \citep{2007Davis}; however, there are two other lines of evidence that might be present in astrometric measurements of globular clusters. |
First. white dwarfs typically have larger proper motions than their progenitors and. older white cdwarls. | First, white dwarfs typically have larger proper motions than their progenitors and older white dwarfs. |
Unfortunately. the errors in the astrometry of voung white chwarts are much larger than their much brighter progenitors on the main sequence. so unless this svstematie is handled carefully. this signal iscillicult to unravel. | Unfortunately, the errors in the astrometry of young white dwarfs are much larger than their much brighter progenitors on the main sequence, so unless this systematic is handled carefully, this signal is difficult to unravel. |
Second. the white cdwarfs after the kick are more likely to lie on racial orbits than their progenitors. | Second, the white dwarfs after the kick are more likely to lie on radial orbits than their progenitors. |
Although measurement errors will weaken observations of this elfect. they cannot mimic an asvnunetry in the orbit parameters. so this signal although weaker is less prone to svstematics and may eventually provide an independent line of evidence for white cdwarl kicks. | Although measurement errors will weaken observations of this effect, they cannot mimic an asymmetry in the orbit parameters, so this signal although weaker is less prone to systematics and may eventually provide an independent line of evidence for white dwarf kicks. |
Furthermore. these calculations indicate that we are on the verge of being able to probe the dynamics of elobular clusters in an entirely new wav through the direct. measurement of the proper motions of the stars. | Furthermore, these calculations indicate that we are on the verge of being able to probe the dynamics of globular clusters in an entirely new way through the direct measurement of the proper motions of the stars. |
Such measurements will provide unique information on the structure of clusters as well as stellar evolution. | Such measurements will provide unique information on the structure of clusters as well as stellar evolution. |
I would like to thank Larvey Richer and Saul Davis for useful discussions. | I would like to thank Harvey Richer and Saul Davis for useful discussions. |
Phe Natural Sciences and Engineering Rescarch Council of Canada. Canadian. Foundation for. Innovation and the British Columbia Ixnowledge. Development Fund supported. this work. | The Natural Sciences and Engineering Research Council of Canada, Canadian Foundation for Innovation and the British Columbia Knowledge Development Fund supported this work. |
Correspondence and. requests. for materials should. be addressed. to hevbephas.ube.ca. | Correspondence and requests for materials should be addressed to heylphas.ubc.ca. |
This research has made use of NASA's Astrophysics Data System Bibliographie Services | This research has made use of NASA's Astrophysics Data System Bibliographic Services |
lines and the strenethoof certain ciission dines. aud all three parameters must vary in correlatioι. | lines and the strength of certain emission lines, and all three parameters must vary in correlation. |
IHowever. observations do not awars support this prediction (sec e.g. Petrov et al. 2001)). | However, observations do not always support this prediction (see e.g. Petrov et al., \cite {petrov01}) ). |
In order to resolve this discrepancy. we sugeest that the area of enhanced Tromospherie ΜΕ ids much more extended than the size of the hot accretio1 spot. | In order to resolve this discrepancy, we suggest that the area of enhanced chromospheric emission is much more extended than the size of the hot accretion spot. |
Iu the models by Romanova et al.. | In the models by Romanova et al., |
less dense iud ixn-uniforii gas surrounds the main accretion stream aud alls onto a much larger area han what is covered by the main stream. up to of the stellar surface. | less dense and non-uniform gas surrounds the main accretion stream and falls onto a much larger area than what is covered by the main stream, up to of the stellar surface. |
A possiülitv is that our observed euliaicect lromosphneriec oenmuüssiou ids excited bv such a mulder ccretion flow and therefore distribute over a sinular area surround18o he siualler hot spot. | A possibility is that our observed enhanced chromospheric emission is excited by such a milder accretion flow and therefore distributed over a similar area surrounding the smaller hot spot. |
One can assuine that the accompanuviug sliocAs aud related phenomena are less pronouiced tiun at tlic Sanam footprints. and that little or no continOUS CXCCSS CLulssjou is released. | One can assume that the accompanying shocks and related phenomena are less pronounced than at the main footprints, and that little or no continuous excess emission is released. |
Du this scenario. a tight ο...ation between the brigifuess. veiling aud CLUISSI1O1 line strength is not expected. | In this scenario, a tight correlation between the brightness, veiling and emission line strength is not expected. |
N-rayv observations of cTTS show that besides verv powerful coroial cluission there is as a rule one conponenut of soft N-ravs that can be associated with the iof spot (seo e.g. Ckdel Telleschi 2007 and Areivofi et al. | X-ray observations of cTTS show that besides very powerful coronal emission there is as a rule one component of soft X-rays that can be associated with the hot spot (see e.g. Güddel Telleschi \cite{gudel07a} and Argiroffi et al. |
2009 aid references therein). | \cite{argiroffi09} and references therein). |
Recently. Drickliotse et al. (2010)) | Recently, Brickhouse et al. \cite{brickhouse10}) ) |
preseuted Xaw spectroscopy of the «Το TW Ίνα, | presented X-ray spectroscopy of the cTTS TW Hya. |
A low-temperature coiiponeut CD — 2.5 \I5) could be identifier with the accretion sliock. as depicted in the stanclare accretion uodels; and a high-temperature colponent was ascribe to coronal enussion a T ~ 10 AUS. | A low-temperature component (T $\sim$ 2.5 MK) could be identified with the accretion shock, as depicted in the standard accretion models, and a high-temperature component was ascribed to coronal emission at T $\sim$ 10 MK. |
In addiion. the authors discuss a third component. represented xo low-eusitv lines of e.g. formed at T ~ 125 NI. which fills a ereat volume outsice the post-shock reeion. | In addition, the authors discuss a third component, represented by low-density lines of e.g. formed at T $\sim$ 1.75 MK, which fills a great volume outside the post-shock region. |
Brickrouse et al. | Brickhouse et al. |
speculate that violet lass outflows and srocks ας propagate from the base of the accreion shock (see also Orlando et al. 2 noy. | speculate that violent mass outflows and shocks may propagate from the base of the accretion shock (see also Orlando et al. \cite{orlando10}) ), |
aud that the pliotosplicvic surroundings are heated. which supplies ionised eas to accretion-fed magnetic structures. such as mmaguctic oops. | and that the photospheric surroundings are heated, which supplies ionised gas to accretion-fed magnetic structures, such as magnetic loops. |
Iu addition. one eau expec that the siuface lavers surrounding the hot spot are excited by UV-lieht directly TOL the hot spot. | In addition, one can expect that the surface layers surrounding the hot spot are excited by UV-light directly from the hot spot. |
Regardless of the «lonunating cause of the eulauced chromospheric enisskn one would expect that the atinosplheric strucure adjacent to the hot spot is affected. | Regardless of the dominating cause of the enhanced chromospheric emission one would expect that the atmospheric structure adjacent to the hot spot is affected. |
lu some carly atteots to explain line ciission and veiling iu cTTS. ITerbiο (1970)) and Cram (1979... 1980)) iutroduced the cowep of a “deep chromosphere”. where the temperature mma occurs at deeper avers in the atinosphere. thus resiting in the appearence of stronger enisson spectrum. | In some early attempts to explain line emission and veiling in cTTS, Herbig \cite{herbig70}) ) and Cram \cite{cram79}, \cite{cram80}) ) introduced the concept of a ”deep chromosphere”, where the temperature minimum occurs at deeper layers in the atmosphere, thus resulting in the appearence of stronger emisson spectrum. |
We speculate that our area of enhanced Cluission ΕΕ he hot spot could be similar to a deep chromosphere with a temperature mininmu much deeper in the atinosprere than i a normal dwarf. but still above the photos)yheric continu. | We speculate that our area of enhanced emission surrounding the hot spot could be similar to a deep chromosphere with a temperature minimum much deeper in the atmosphere than in a normal dwarf, but still above the photospheric continuum. |
Iu anv case. thle chromospheric emissio rin cTTS appears to be related not oulv to solar-like macetic activity. but is fo a considerable exteut powered by accretion processes. | In any case, the chromospheric emission in cTTS appears to be related not only to solar-like magnetic activity, but is to a considerable extent powered by accretion processes. |
We have stucied periodic changes iu photospheric lies and narrow components of emission lines in DR Tauri aud other classical T Tauri stars. aud investigated the nature of the dilution of plotospheric lines. the so-called veiling. | We have studied periodic changes in photospheric lines and narrow components of emission lines in DR Tauri and other classical T Tauri stars, and investigated the nature of the dilution of photospheric lines, the so-called veiling. |
Iu particular we note: We suggest that the area of cnhanced chromospheric chussion comes from an extended region at the stellar surface surrounding the hot spot. | In particular we note: We suggest that the area of enhanced chromospheric emission comes from an extended region at the stellar surface surrounding the hot spot. |
The plivsical mechanisin hat trigecrs the enissionu cau be related to current ideas hat differ from the standard model in that accretion Is nore wide-spread. or cluanates from injection of uass and energv from the shocked region under the naib accretion stream. | The physical mechanism that triggers the emission can be related to current ideas that differ from the standard model in that accretion is more wide-spread, or emanates from injection of mass and energy from the shocked region under the main accretion stream. |
As a result the photosphere surrounding the hot spot is leaed. which leads to a uodified. atmospheric structure. which iu tum produces he chhanced cussion n narro emission lines responsible or the extra compoucut in the veiling through line-filliug of photospleric absorption lines. | As a result the photosphere surrounding the hot spot is heated, which leads to a modified atmospheric structure, which in turn produces the enhanced emission in narrow emission lines responsible for the extra component in the veiling through line-filling of photospheric absorption lines. |
This ts the “original” Tonry calibration. equation (1)). applied to all This is equation (1)) applied for (V—Dg>1.09. and the following equation (2)) — which has half the slope of equation (1)) - applied for (V—Όος1.09: This case represents the theoretical SBF models plottec in the left panel of Fig. 1. | This is the “original” Tonry calibration, equation \ref{sbfrel}) ), applied to all This is equation \ref{sbfrel}) ) applied for $(V-I)_0>1.09$, and the following equation \ref{sbfrelmod}) ) – which has half the slope of equation \ref{sbfrel}) ) – applied for $(V-I)_0<1.09$: This case represents the theoretical SBF models plotted in the left panel of Fig. \ref{twobranch}, |
which viewed as a whole suggest a non-zero but shallower slope than that of the empirical calibration at redder colours. | which viewed as a whole suggest a non-zero but shallower slope than that of the empirical calibration at redder colours. |
See case C) below for the reasoi of choosing (V—7)=1.09 as colour This is equation (1)) applied for (V—Do>1.09 anc a constant of M;=—2.00 mag applied for (V—Da<1.09. | See case C) below for the reason of choosing $(V-I)=1.09$ as colour This is equation \ref{sbfrel}) ) applied for $(V-I)_0>1.09$ and a constant of $\overline{M}_I=-2.00$ mag applied for $(V-I)_0<1.09$. |
This represents the theoretical SBF models plotted in the right panel of Fig. 1.. | This represents the theoretical SBF models plotted in the right panel of Fig. \ref{twobranch}. |
The specific limit of(V—7)5=1.09 is choser because at this colour. the colour independent M,~—2.0 mag of the models in the right panel of Fig. | The specific limit of $(V-I)_0=1.09$ is chosen because at this colour, the colour independent $\overline{M}_I\simeq -2.0$ mag of the models in the right panel of Fig. |
|) matches the value of M; from equation (1). | \ref{twobranch} matches the value of $\overline{M}_I$ from equation \ref{sbfrel}) ). |
This is also the reason for adopting the same limit in case B). since the latter case represents the "hybrid" solution between case A) and case A flat and a steep calibration relation are. valid in the same colour regime. | This is also the reason for adopting the same limit in case B), since the latter case represents the “hybrid” solution between case A) and case A flat and a steep calibration relation are valid in the same colour regime. |
Depending on the strength. of the SBF signal. one of these is applied to a given galaxy. | Depending on the strength of the SBF signal, one of these is applied to a given galaxy. |
This case represents the "two branch" hypothesis put forward in the investigations. by Jerjen et al. | This case represents the “two branch” hypothesis put forward in the investigations by Jerjen et al. |
The steep relation is equation (1)). | The steep relation is equation \ref{sbfrel}) ). |
It is applied to galaxies which have (VW-—fo>1.09or whose distance in case B) is smaller than the mean distance. | It is applied to galaxies which have $(V-I)_0>1.09$ whose distance in case B) is smaller than the mean distance. |
The flat relation is simply a constant M,=--00 mag. like in case C). | The flat relation is simply a constant $\overline{M}_I=-2.00$ mag, like in case C). |
This is applied to all galaxies with (V—Do<1.09 whose distance in case B) is larger than the mean distance. | This is applied to all galaxies with $(V-I)_0<1.09$ whose distance in case B) is larger than the mean distance. |
In other words. galaxies with 77; above the line for ease B) in Fig. | In other words, galaxies with $\overline{m}_I$ above the line for case B) in Fig. |
|. are assigned M;=—2.00 mag. while those below are assigned relation (1)). | \ref{twobranch} are assigned $\overline{M}_I=-2.00$ mag, while those below are assigned relation \ref{sbfrel}) ). |
The mean distance obtained with this case ts of course very close to that in case B). | The mean distance obtained with this case is of course very close to that in case B). |
Those four cases will be tested against the real SBF measurements in Sect. 4.. | Those four cases will be tested against the real SBF measurements in Sect. \ref{testsbf}. |
The observations for this paper were taken in the nights 25th of October 2003 and 5-7 of December 2004 at Las Campanas Observatory (LCO). Chile. | The observations for this paper were taken in the nights 25th of October 2003 and 5-7 of December 2004 at Las Campanas Observatory (LCO), Chile. |
The instrument used was the Inamori Magellan Areal Camera and Spectrograph “IMACS” in imaging mode with the “short” f/2 camera. mounted at the 6.5m Baade telescope. | The instrument used was the Inamori Magellan Areal Camera and Spectrograph “IMACS” in imaging mode with the “short” f/2 camera, mounted at the 6.5m Baade telescope. |
The double-asphere. glass-and-oil-Iens f/2 camera produces an image of 27.4 field diameter at 0.20 aresee per pixel. | The double-asphere, glass-and-oil-lens f/2 camera produces an image of $'$ field diameter at 0.20 arcsec per pixel. |
The field is vignetted in the corners (by the tertiary mirror and its mounting assembly). changing from flux loss at R = 12’ to. at R = 15’. | The field is vignetted in the corners (by the tertiary mirror and its mounting assembly), changing from flux loss at R = $'$ to at R = $'$. |
The IMACS detector is an 8192x8192 CCD mosaic camera which uses 8 thinned 2kx4k detectors. | The IMACS detector is an $\times$ 8192 CCD mosaic camera which uses 8 thinned $\times$ 4k detectors. |
Gaps of about 50 pixels separate the In total. 7 fields in the central Fornax cluster were observed in the two bands V and /. | Gaps of about 50 pixels separate the In total, 7 fields in the central Fornax cluster were observed in the two bands $V$ and $I$. |
The total integration times in V were between 1200 and 1800 seconds except for the central field around NGC 1399, for which it was 900 seconds. | The total integration times in $V$ were between 1200 and 1800 seconds except for the central field around NGC 1399, for which it was 900 seconds. |
The total /-band integration time was between 3800 and 5160 seconds. except for the central NGC 1399 field. which had 1800 seconds Standard star images of the Landolt fields SA92. SA9S. SA98 and SAI04 -- the latter one only in the second rur — Were taken at different airmass values in one short (2 seconds) and one long exposure (10 seconds). | The total $I$ -band integration time was between 3800 and 5160 seconds, except for the central NGC 1399 field, which had 1800 seconds Standard star images of the Landolt fields SA92, SA95, SA98 and SA104 – the latter one only in the second run – were taken at different airmass values in one short (2 seconds) and one long exposure (10 seconds). |
This resulted in about 110 data points for the first run and about 140 for the second one. each time corresponding to about 40 standard stars. | This resulted in about 110 data points for the first run and about 140 for the second one, each time corresponding to about 40 standard stars. |
Comparison of calibrated colours and magnitudes for objects imaged in both runs yielded absolute zero point uncertainties of the order 0.02 mag. | Comparison of calibrated colours and magnitudes for objects imaged in both runs yielded absolute zero point uncertainties of the order 0.02 mag. |
The image reduction steps before the SBF measurement were the following: first. a master-bias was created for each chip and was subtracted from the domeflat exposures. | The image reduction steps before the SBF measurement were the following: first, a master-bias was created for each chip and was subtracted from the domeflat exposures. |
Then for each chip the bias corrected dome flats were combined to master-domeflats. | Then for each chip the bias corrected dome flats were combined to master-domeflats. |
Having the master-biases and master-domeflats prepared for each chip. we used the COSMOS to do bias-subtraction. trimming and flat-field correction of the raw science frames. | Having the master-biases and master-domeflats prepared for each chip, we used the COSMOS to do bias-subtraction, trimming and flat-field correction of the raw science frames. |
The reduced single science frames were registered with integer pixel shifts to avoid distortior of the SBF power spectrum and combined using a min-max rejection algorithm. | The reduced single science frames were registered with integer pixel shifts to avoid distortion of the SBF power spectrum and combined using a min-max rejection algorithm. |
The seeing FWHM of the combined images generally ranged between 0.6 and 0.9". except at the edges of the field of view: especially for the first run i 2003. when the atmospheric dispersion corrector had not beer installed yet. distortions close to the image border caused seeing degradations to about 2". | The seeing FWHM of the combined images generally ranged between 0.6 and $''$, except at the edges of the field of view: especially for the first run in 2003, when the atmospheric dispersion corrector had not been installed yet, distortions close to the image border caused seeing degradations to about $''$. |
There were two Fornax dEs for which we could not measure SBF because they were located in these regions of high The same image reduction. procedure. às for the single science frames was performed for the standard star images. | There were two Fornax dEs for which we could not measure SBF because they were located in these regions of high The same image reduction procedure as for the single science frames was performed for the standard star images. |
Their instrumental magnitudes were measured with the IRAF package APPHOT in apertures equal to those used by Landolt (1994)). | Their instrumental magnitudes were measured with the IRAF package APPHOT in apertures equal to those used by Landolt \cite{Landolxx}) ). |
Then. a single photometric solution was determined for the entire 8kx8k image. | Then, a single photometric solution was determined for the entire $\times$ 8k image. |
Chip-to-chip variations around the mean photometric zero point were smaller than 0.02 mag. | Chip-to-chip variations around the mean photometric zero point were smaller than 0.02 mag. |
To correct for galactic reddening and absorption. we used the values from Schlegel et al. (1998)). | To correct for galactic reddening and absorption, we used the values from Schlegel et al. \cite{Schleg98}) ), |
who give A;=0.025 and £(V—7)=0.018 for the coordinates of the Fornax cluster. | who give $A_I=0.025$ and $E(V-I)=0.018$ for the coordinates of the Fornax cluster. |
The SBF measurement procedure was similar to that already outlined 1n. Mieske et al. (2005)) | The SBF measurement procedure was similar to that already outlined in Mieske et al. \cite{Mieske05}) ) |
based on VLT FORSI photometry. | based on VLT FORS1 photometry. |
Some nuances have been modified: the fluctuations from the sky background were now derived | Some nuances have been modified: the fluctuations from the sky background were now derived |
The matter distribution of the NEW modelNavarroetal.(1907). is given by where p.=BIT?(Απ) is the critical clensity of the universe as a function of redshift. | The matter distribution of the NFW model\cite{NFW} is given by where $\rho_c = 3 H^2(z)/(8\pi G)$ is the critical density of the universe as a function of redshift. |
The two paralcters of the model are the scale radius. rs. and à... In the literature à, is often replaced by the concentration parameter e defined by and the scale radius rj is replaced by the virial radius rogo=ery. Which is the radius inside which the average mass density is 200p... | The two parameters of the model are the scale radius, $r_s$, and $\delta_c$ In the literature $\delta_c$ is often replaced by the concentration parameter $c$ defined by and the scale radius $r_s$ is replaced by the virial radius $r_{200} = c \, r_s$, which is the radius inside which the average mass density is $200 \rho_c$. |
Like the SIS model. the total mass over all space is undefined. so we iust truncate the model at sole radius rj. | Like the SIS model, the total mass over all space is undefined, so we must truncate the model at some radius $r_c$. |
However. because the NEW model is indexed to p.. the mass interior to a coustaut proper radius varies with redshift. | However, because the NFW model is indexed to $\rho_c$, the mass interior to a constant proper radius varies with redshift. |
We choose to hold the mass of the halo coustaut and truuncate the NEW profile at rogo. so that the constant total mass of the halo is The truncation radius r,=reyy is then a known function of redshift through p: Iu practice. we pick a value for rogo at the time the lightray passes the lens aud use Eq. | We choose to hold the mass of the halo constant and truncate the NFW profile at $r_{200}$, so that the constant total mass of the halo is The truncation radius $r_c = r_{200}$ is then a known function of redshift through $\rho_c$: In practice, we pick a value for $r_{200}$ at the time the lightray passes the lens and use Eq. |
321. to compute constant halo mass and Eq. | \ref{M200} to compute constant halo mass and Eq. |
32. to determine the truncation radius as a function of time. | \ref{r200c} to determine the truncation radius as a function of time. |
For. umrogo <1. the euclosed nass is where eis the concentration parameter. | For $x \equiv r_p/r_{200} < 1$ , the enclosed mass is where $c$ is the concentration parameter. |
The NEW eravitational poteutial is obtained by sinoothlv matching the spherically svuuuctric eravitational poteutial for the mass density elven by Eq. | The NFW gravitational potential is obtained by smoothly matching the spherically symmetric gravitational potential for the mass density given by Eq. |
29 to a point mass potential at «©=1. | \ref{rho_nfw} to a point mass potential at $x=1$. |
The result is where for cosmological spacetimes. we take 6=a(f)irírogg. With rogo defined in Eq. | The result is where for cosmological spacetimes, we take $x = a(t)
r/r_{200}$ , with $r_{200}$ defined in Eq. |
32. and constaut concentration parameter c. | \ref{r200c} and constant concentration parameter $c$. |
We begin our discussionof NEW. thin-lens models by determining the two-dimensional. projected lnass densitv used in eravitational lensing with an arbitrary truncation radius r.. | We begin our discussionof NFW thin-lens models by determining the two-dimensional, projected mass density used in gravitational lensing with an arbitrary truncation radius $r_c$ . |
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