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They invoke two heating mechanisms: the magneticwave and the acousticwave heating. | They invoke two heating mechanisms: the magnetic–wave and the acoustic–wave heating. |
Theoretical chromospheric fluxes for the former the observed fux of the active stars. while theoretical fluxes for the latter mechauisia match the observed fiux of the inactive stars. | Theoretical chromospheric fluxes for the former mechanism match the observed flux of the active stars, while theoretical fluxes for the latter mechanism match the observed flux of the inactive stars. |
In order to provide a physical expanation for the occurrence of the VP eap. Durnevctal.(1981) proposed a transition from a coniplex to a siupler magneticfed inorphology which occurs at he time when the rotation decreases choug[um to reach a Έποςtold value. | In order to provide a physical explanation for the occurrence of the VP gap, \cite{dmr81} proposed a transition from a complex to a simpler magnetic–field morphology which occurs at the time when the rotation decreases enough to reach a threshold value. |
More recently. Barnes(2003 detected fwo sequences iu the 21οςversuscolour diagram oope 1 clusters. aud le associated them with two different rotation morpholoeies, intertwined with stellar magnetic fields. | More recently, \cite{barnes} detected two sequences in the period–versus–colour diagram of open clusters, and he associated them with two different rotation morphologies, intertwined with stellar magnetic fields. |
Bolun-Viteise(2007) suggested a change of diano mechauisui to explain the fact that stars occupy two very distinct secuences In a rotatio- periodversuscycle period diagram. | \cite{bv07} suggested a change of dynamo mechanism to explain the fact that stars occupy two very distinct sequences in a rotation period–versus–cycle period diagram. |
Not all of these works can be used to explain our results. | Not all of these works can be used to explain our results. |
What is relevant for the presceut discussion is the possibility of a change in the nature of the dvuau10 1uechauisui talking place at a specific point of the mainsequence lifetime of solartype stars. | What is relevant for the present discussion is the possibility of a change in the nature of the dynamo mechanism taking place at a specific point of the main–sequence lifetime of solar–type stars. |
The uost important conclusion of our analysis is that. before aud after the fast drop of CA from above to beow the VP eap in only 200 M3r. its alleged simootl decay must be nuch less iu)ortant than short term variations. | The most important conclusion of our analysis is that, before and after the fast drop of CA from above to below the VP gap in only 200 Myr, its alleged smooth decay must be much less important than short term variations. |
We noted first in Pace&Pasquini(2001) that οA iu intermediate age clusters had already droppce to the solar evel. | We noted first in \cite{paper1} that CA in intermediate age clusters had already dropped to the solar level. |
Àter the reanalysis of the a using temperatre decrimined spectroscopicallv from ion lines. 1 out of the 7 stars in the two crinediate aee clusters IC 1651 and NCC 3680 turned. Ol to liein the VP ea» | After the reanalysis of the data using temperature determined spectroscopically from iron lines, 1 out of the 7 stars in the two intermediate age clusters IC 4651 and NGC 3680 turned out to lie in the VP gap. |
Iu addition. the membership of one of the two stars i1 NGC 3680 has been questioned Aithom-Twarogeta. | In addition, the membership of one of the two stars in NGC 3680 has been questioned \cite{chato}. |
(2009).. IILowever. the data still point uudoubtedlv to ἃ lack of evolution after 1.1 Cir. slwe 5 out of 6 secwe Intermediate age stars have CA levels equal or ower to those spanned by the Sun and by M 67 sars (clupty cdots in Figure 1)). | However, the data still point undoubtedly to a lack of evolution after 1.4 Gyr, since 5 out of 6 secure intermediate age stars have CA levels equal or lower to those spanned by the Sun and by M 67 stars (empty dots in Figure \ref{fig}) ). |
Tus conclusion is also corroborated bv the work of Lyra&PortodeMelo(2005). | This conclusion is also corroborated by the work of \cite{lyra}. |
. The other part of the conclusions. that regarding the time interval between 0.7 and 1.2 Cr (filles dots in Figure 1)). was achieved after the analvsis of 5 stars In NGC 5822 and IC £756. together with the older data ou IIvades aid Praesepe (Paceeta.2009). | The other part of the conclusions, that regarding the time interval between 0.7 and 1.2 Gyr (filled dots in Figure \ref{fig}) ), was achieved after the analysis of 5 stars in NGC 5822 and IC 4756, together with the older data on Hyades and Praesepe \citep{letter}. |
.. We could show tha metallicity is uulikelv to play a niajor role. since there is no siguificaut cüffereuce )etween three almost coeval oon clusters at different imetallicities: Pracsepe. a Fe/II|2t.27 (Pacectal.2008).. ITvades. at |Fe/TI|-0.13 al. 2003).. and IC 1756 at [Fe/TJ=0.d (Paceot 2010).. | We could show that metallicity is unlikely to play a major role, since there is no significant difference between three almost coeval open clusters at different metallicities: Praesepe, at [Fe/H]=0.27 \citep{papoc1}, Hyades, at [Fe/H]=0.13 \citep{paulson}, and IC 4756 at [Fe/H]=0.01 \citep{papoc2}. . |
The quoted iron abundauce or Pracsepe is somewhat differeut from other published works (seee.g. 2009).. | The quoted iron abundance for Praesepe is somewhat different from other published works \cite[see
e.g.][]{chato}. . |
ILowever. this result is mechanisimfpased ou the highquality spectra deseribed in Section | However, this result is based on the high–quality spectra described in Section |
luminosity ratios for the Compton thick AGN NGC 4945 (squares) and Circinus (diamonds) respectively, calculated using both the observed (left point) and (right point) 2-10keV luminosity. | luminosity ratios for the Compton thick AGN NGC 4945 (squares) and Circinus (diamonds) respectively, calculated using both the observed (left point) and (right point) 2-10keV luminosity. |
From Fig. 3,, | From Fig. \ref{fig:irx12}, |
Group 1 AGN (black circles) are the more luminous AGN and have a remarkably narrow dispersion in luminosity ratio (30«γα« 1) for most (97/107) Group 1 AGN. | Group 1 AGN (black circles) are the more luminous AGN and have a remarkably narrow dispersion in luminosity ratio $30<R_{ir/x}<1$ ) for most (97/107) Group 1 AGN. |
Of the ten Group 1 AGN with &;,/;>30, four are cross-classified (i.e. Seyfert 1 plus starburst) and four are heavily obscured Seyfert 1s (Ng10cm ?). | Of the ten Group 1 AGN with $R_{ir/x}>30$, four are cross-classified (i.e. Seyfert 1 plus starburst) and four are heavily obscured Seyfert 1s $N_{H} \sim 10^{23} \rm{cm}^{-2}$ ). |
So Group 1 AGN with >30 could be accounted for either by the addition of IR [οluminosity from star formation or a drop in observed X-ray luminosity due to a change in partial-covering absorption (McKernan&Yaqoob1998;Risaliti 2008).. | So Group 1 AGN with $R_{ir/x}>30$ could be accounted for either by the addition of IR luminosity from star formation or a drop in observed X-ray luminosity due to a change in partial-covering absorption \citep{b91,b92,b93}. |
So even with our simple approach, Fig. | So even with our simple approach, Fig. |
3 indicates that ~90% of Group 1 AGN in our sample have 1«R;,/,30. | \ref{fig:irx12} indicates that $\sim 90\%$ of Group 1 AGN in our sample have $1<R_{ir/x}<30$. |
By contrast, Group 2 AGN (red crosses) generally have lower X-ray luminosities and have a wider dispersion in R,,/, than Group 1 AGN. | By contrast, Group 2 AGN (red crosses) generally have lower X-ray luminosities and have a wider dispersion in $R_{ir/x}$ than Group 1 AGN. |
In contrast to the Group 1 AGN, 97/109 group 2 AGN have Rir/e>30, with 83/109 group 2 AGN having Ίνα>100. | In contrast to the Group 1 AGN, 97/109 group 2 AGN have $R_{ir/x}>30$, with 83/109 group 2 AGN having $R_{ir/x}>100$. |
Of the 12/109 Group 2 AGN with 1<R30 (ie. overlapping with the main Group 1 population), we find that six have associated radio jets and two are Compton-thin Seyfert 2 AGN. | Of the 12/109 Group 2 AGN with $1<R<30$ (i.e. overlapping with the main Group 1 population), we find that six have associated radio jets and two are Compton-thin Seyfert 2 AGN. |
From Table 2,, 19/221 AGN in our sample are non-Seyferts, of which 10 are LINERs, 8 LINERs with Hy regions and 1 QSO. | From Table \ref{tab:z}, 19/221 AGN in our sample are non-Seyferts, of which 10 are LINERs, 8 LINERs with $H_{\rm{II}}$ regions and 1 QSO. |
The LINERs lie in the bottom left hand corner of the Group 2 distribution (Lir<0.5x10% erg/s,L,<10°° erg/s) and the LINERs with Hy regions span a slightly larger range of the Group 2 distribution (Li,<5x10” erg/s,L,«10? erg/s). | The LINERs lie in the bottom left hand corner of the Group 2 distribution $L_{ir}<0.5 \times
10^{42}$ $L_{x}<10^{39}$ erg/s) and the LINERs with $H_{\rm{II}}$ regions span a slightly larger range of the Group 2 distribution $L_{ir}<5
\times 10^{42}$ $L_{x}<10^{40}$ erg/s). |
A key point from Fig. | A key point from Fig. |
3 is that the luminosity ratio for NGC 4945 and Circinus as calculated from the interpolated intrinsic 2-10keV luminosity (right point) lies in the main luminosity ratio band of Group 1 AGN. | \ref{fig:irx12} is that the luminosity ratio for NGC 4945 and Circinus as calculated from the interpolated intrinsic 2-10keV luminosity (right point) lies in the main luminosity ratio band of Group 1 AGN. |
This is a nice demonstration of the fact that behind the Compton-thick obscuration of NGC 4945 and Circinus, lies the same phenomenon as in most Group 1 AGN. | This is a nice demonstration of the fact that behind the Compton-thick obscuration of NGC 4945 and Circinus, lies the same phenomenon as in most Group 1 AGN. |
Without very hard X-ray measurements, it is very difficult to estimate intrinsic 2-10keV X-ray luminosities for most Group 2 AGN in our sample and certainly any such estimates would be highly model-dependent. | Without very hard X-ray measurements, it is very difficult to estimate intrinsic 2-10keV X-ray luminosities for most Group 2 AGN in our sample and certainly any such estimates would be highly model-dependent. |
Nevertheless, the examples of NGC 4945 and Circinus indicate that the luminosity ratios for many Group 2 AGN may be similar to the Group 1 luminosity ratio range (~30<Ri,/v& 1). | Nevertheless, the examples of NGC 4945 and Circinus indicate that the luminosity ratios for many Group 2 AGN may be similar to the Group 1 luminosity ratio range $\sim 30 \leq R_{ir/x} \leq 1$ ). |
However, there are many model-dependent caveats in estimating the intrinsic X-ray luminosity of an AGN, so in refsec:discuss below we will take a different approach to Lutzetal.(2001);KrabbeHorst and attempt to derive constraints on AGN using luminosities. | However, there are many model-dependent caveats in estimating the intrinsic X-ray luminosity of an AGN, so in \\ref{sec:discuss} below we will take a different approach to \citet{b3,b4,b5} and attempt to derive constraints on AGN using luminosities. |
Figure 4 shows the mean far-IR (100um) luminosity of the AGN in our sample versus the mean 2-10keV observed X-ray luminosity. | Figure \ref{fig:irx100} shows the mean far-IR $\micron$ ) luminosity of the AGN in our sample versus the mean 2-10keV observed X-ray luminosity. |
Note that several AGN had NED flux measurements at 124m but only upper limits at 100um and vice versa, so there is not a perfect one-to-one correspondence between all points on Fig. | Note that several AGN had NED flux measurements at $12\micron$ but only upper limits at $100\micron$ and vice versa, so there is not a perfect one-to-one correspondence between all points on Fig. |
4 and Fig. 3.. | \ref{fig:irx100} and Fig. \ref{fig:irx12}. |
Nevertheless, the dispersion in for the Group 2 AGN in particular seems to be larger in | Nevertheless, the dispersion in $R_{ir/x}$ for the Group 2 AGN in particular seems to be larger in Fig. |
R;,/,Fig. 4 than in Fig. | \ref{fig:irx100} than in Fig. |
3 and the various archetypal Group 2 AGN appear more separated in luminosity ratio. | \ref{fig:irx12} and the various archetypal Group 2 AGN appear more separated in luminosity ratio. |
For example, the starburst galaxy Arp 220 (open star) is much more clearly distinguished from the classic Type 2 AGN NGC 1068 (open triangle) in Fig. | For example, the starburst galaxy Arp 220 (open star) is much more clearly distinguished from the classic Type 2 AGN NGC 1068 (open triangle) in Fig. |
4 than in Fig 3.. | \ref{fig:irx100} than in Fig \ref{fig:irx12}. . |
The non-Seyfert AGN, e.g. the LINERs and LINERs with Hr: regions, are also highlighted. | The non-Seyfert AGN, e.g. the LINERs and LINERs with $\rm{H}_{II}$ regions, are also highlighted. |
As we should expect from such low X-ray and IR luminosity sources, LINERs occupy the lower left-hand corner of the Group 2 AGN dispersion in Fig. 4.. | As we should expect from such low X-ray and IR luminosity sources, LINERs occupy the lower left-hand corner of the Group 2 AGN dispersion in Fig. \ref{fig:irx100}. |
The LINERs with associated Hr regions have generally higher IR luminosities than the 'pure' LINERs, but the X-ray luminosities remain fairly low. | The LINERs with associated $\rm{H}_{II}$ regions have generally higher IR luminosities than the 'pure' LINERs, but the X-ray luminosities remain fairly low. |
Figures 3 and 4 both show a very clear distinction between the observed luminosity ratios (and their dispersions) for group 1 and group 2 AGN. | Figures \ref{fig:irx12} and \ref{fig:irx100} both show a very clear distinction between the observed luminosity ratios (and their dispersions) for group 1 and group 2 AGN. |
However, based on the modelled intrinsic 2-10keV X-ray luminosity, Krabbeetal.(2001);Horst(2007) found that the IR to X-ray luminosity ratios for group 1 and group 2 AGN were very similar. | However, based on the modelled intrinsic 2-10keV X-ray luminosity, \citet{b4,b5} found that the near-IR to X-ray luminosity ratios for group 1 and group 2 AGN were very similar. |
Figure 5 is as Fig. 3,, | Figure \ref{fig:lutz} is as Fig. \ref{fig:irx12}, |
but with open circles denoting AGN from the study by Horstetal.(2007) and the ranges of luminosity ratio established by Horstetal.(2007) for group 1 (solid lines) and group 2 (dashed lines) AGN. | but with open circles denoting AGN from the study by \citet{b5} and the ranges of luminosity ratio established by \citet{b5} for group 1 (solid lines) and group 2 (dashed lines) AGN. |
The dispersion in the luminosity ratio of our group 1 AGN matches that found by Horstetal.(2007) quite well, indicating that observed and intrinsic 2-10keV X-ray luminosities are similar in most group 1 AGN. | The dispersion in the luminosity ratio of our group 1 AGN matches that found by \citet{b5} quite well, indicating that observed and intrinsic 2-10keV X-ray luminosities are similar in most group 1 AGN. |
Our group 2 population seems to diverge dramatically in luminosity ratio from the range found by Horstetal.(2007),, suggesting that the observed and intrinsic 2-10keV luminosities in group 2 AGN are indeed dramatically different. | Our group 2 population seems to diverge dramatically in luminosity ratio from the range found by \citet{b5}, suggesting that the observed and intrinsic 2-10keV luminosities in group 2 AGN are indeed dramatically different. |
However, Horstet estimate the intrinsic X-ray luminosity in the group 2 AGN based on different model fits in the literature, which (a) can vary for the same AGN dataset and which (b) lead tomodel-dependent constraints on AGN structure (e.g. a clumpy torus as suggested by Horstetal. (2007))). | However, \citet{b5} estimate the intrinsic X-ray luminosity in the group 2 AGN based on different model fits in the literature, which (a) can vary for the same AGN dataset and which (b) lead tomodel-dependent constraints on AGN structure (e.g. a clumpy torus as suggested by \citet{b5}) ). |
By contrast, an IR-X-ray luminosity ratio based on the | By contrast, an IR-X-ray luminosity ratio based on the |
respect to the one used by Lazzati et al. ( | respect to the one used by Lazzati et al. ( |
2001). depending on the considered band. by ~0.1-0.2 mag. | 2001), depending on the considered band, by $\sim$ 0.1-0.2 mag. |
We have evaluated this discrepancy by comparing our and their quasi-simultaneous («2 hr) measurements of September 28. 2000. and have corrected their photometry to obtain a mutually consistent zero-point level. | We have evaluated this discrepancy by comparing our and their quasi-simultaneous $<$ 2 hr) measurements of September 28, 2000, and have corrected their photometry to obtain a mutually consistent zero-point level. |
Because of the above procedure. we did not plot in Fig. | Because of the above procedure, we did not plot in Fig. |
2 the measurements acquired by Lazzati et al. ( | 2 the measurements acquired by Lazzati et al. ( |
2001) on that date. | 2001) on that date. |
It should also be noted that the photometry errors quoted throughout the rest of the paper are only statistical and do not account for any possible zero-point offset. which we expect to be smaller than 2%. | It should also be noted that the photometry errors quoted throughout the rest of the paper are only statistical and do not account for any possible zero-point offset, which we expect to be smaller than . |
. NIR imaging in J. H and Κι bands was obtained between September 21 and December 22. 2000. at the ESO NTT+Sofl in La Silla (Chile). and again at the in Paranal (Chile). equipped with ISAAC (see Table 1). | NIR imaging in $J$, $H$ and $K_s$ bands was obtained between September 21 and December 22, 2000, at the ESO NTT+SofI in La Silla (Chile), and again at the in Paranal (Chile), equipped with ISAAC (see Table 1). |
The Sofl NIR camera carried a Rockwell Hawai 1024x1024 pixel HeCdTe array for imaging and spectroscopy in the 0.9-2.5 pm band. | The SofI NIR camera carried a Rockwell Hawaii $\times$ 1024 pixel HgCdTe array for imaging and spectroscopy in the 0.9–2.5 $\mu$ m band. |
The plate scale was 07292 pix”! and the corresponding field of view was 4f9x4'9. | The plate scale was $\farcs$ 292 $^{-1}$ and the corresponding field of view was $\farcm$ $\times$ $\farcm$ 9. |
ISAAC was equipped. in the short-wavelength (0.9—2.5 jm) NIR range. with à similar Rockwell Hawan 1024x1024 pixel HeCdTe array which had a scale of OV 148 pix”! and a field of view of 2/5x2/5. | ISAAC was equipped, in the short-wavelength (0.9–2.5 $\mu$ m) NIR range, with a similar Rockwell Hawaii $\times$ 1024 pixel HgCdTe array which had a scale of $\farcs$ 148 $^{-1}$ and a field of view of $\farcm$ $\times$ $\farcm$ 5. |
The K, filter is centered at 2.12 gm and has a FWHM of 0.34 um. For each NIR observation the total integration time was split into images of 40 s each. with dithering after each individual exposure. | The $K_s$ filter is centered at 2.12 $\mu$ m and has a FWHM of 0.34 $\mu$ m. For each NIR observation the total integration time was split into images of 40 s each, with dithering after each individual exposure. |
Reduction of the NIR images was performed with the IRAF and STSDASpackages’. | Reduction of the NIR images was performed with the IRAF and STSDAS. |
.. Each image was reduced by first subtracting a mean sky. obtained from the median of a number of frames acquired just before and after each processed image. | Each image was reduced by first subtracting a mean sky, obtained from the median of a number of frames acquired just before and after each processed image. |
Before frames were used for sky subtraction. stars in them were eliminated by à background interpolation algorithm (imedit) combined with an automatic "star finder” (daofind). | Before frames were used for sky subtraction, stars in them were eliminated by a background interpolation algorithm ) combined with an automatic “star finder” ). |
Then. a differential dome flattield correction was applied to the sky-subtracted image. and the frames were registered to fractional pixels and combined. | Then, a differential dome flatfield correction was applied to the sky-subtracted image, and the frames were registered to fractional pixels and combined. |
The telescope dithering was measured from the offsets of field objects in each image and the images were averaged together using inter-pixel shifts. | The telescope dithering was measured from the offsets of field objects in each image and the images were averaged together using inter-pixel shifts. |
Magnitudes were measured inside circular apertures of diameter comparable tothe FWHM of each image. | Magnitudes were measured inside circular apertures of diameter comparable tothe FWHM of each image. |
We calibrated the NIR photometry with stars selected from the NICMOS Standards List (Persson et al. | We calibrated the NIR photometry with stars selected from the NICMOS Standards List (Persson et al. |
1998). | 1998). |
The | The |
and more microscopic models. | and more microscopic models. |
Lhe range we Line is 0.05 - OFae ye"—7.10"nLO14 ο δν | The range we find is 0.05 - 0.07 $^{-3} \equiv 7 \times 10^{13} - 10^{14}$ g $^{-3}$. |
"Phe cillerences in our methods makes a direct. comparison of our results cillicult. | The differences in our methods makes a direct comparison of our results difficult. |
For example. Sotani(2011). obtains a fundamental moce frequeney of z22 Hz for a L4Al. neutron star when the shear modulus is taken to be that of a Coulomb crystal up to the crust-core boundary. compared. with z34 Iz in our moclel. taken at £=46 MeV: however. the stiffness of the crust in the Sotani model would lead to a thinner crust and a decrease in the frequency. compared: with our results. | For example, \citet{Sotani2011} obtains a fundamental mode frequency of $\approx 22$ Hz for a $1.4 M_{\odot}$ neutron star when the shear modulus is taken to be that of a Coulomb crystal up to the crust-core boundary, compared with $\approx 34$ Hz in our model, taken at $L = 46$ MeV; however, the stiffness of the crust in the Sotani model would lead to a thinner crust and a decrease in the frequency compared with our results. |
Lt would be useful to make a comparison of the methods used. using consistent nuclear physies. | It would be useful to make a comparison of the methods used using consistent nuclear physics. |
As a final note. we have estimated the comparative ellect of the pasta phases and the effects. of superfluid neutrons. finding them to be comparable. as suggested in Sotani(2011). | As a final note, we have estimated the comparative effect of the pasta phases and the effects of superfluid neutrons, finding them to be comparable, as suggested in \citet{Sotani2011}. |
This work is supported in part by the National jJXeronauties and Space Xdministration. uncer grant ΑΝΝΑ] issued. through the Science. Mission. Directorate. and. the National Science Foundation under grants. PIIY-0757839 and PIIY-1068022 and the Texas Coordinating Board of Iligher Exlucation under grant No. | This work is supported in part by the National Aeronautics and Space Administration under grant NNX11AC41G issued through the Science Mission Directorate and the National Science Foundation under grants PHY-0757839 and PHY-1068022 and the Texas Coordinating Board of Higher Education under grant No. |
003565-0004-2007. | 003565-0004-2007. |
range of intermecliate redshifts. | range of intermediate redshifts. |
? find a similar result which suggests stronger evolution in the redshift number density of strong aabsorption svstems relative to lower LEW samples. albeit over a smaller redshift range 0.4xz:0.5. | \citet{2009ApJ...698..819L} find a similar result which suggests stronger evolution in the redshift number density of strong absorption systems relative to lower EW samples, albeit over a smaller redshift range $0.4 \leq z \leq 0.8$. |
Using our parametrized dust correction (equation 7)). we predict that the relative fraction of high. EW systems will decrease. by pper cent over the redshift interval ΟΕ<z2.0. | Using our parametrized dust correction (equation \ref{eq:fullebvew}) ), we predict that the relative fraction of high EW systems will decrease by per cent over the redshift interval $0.4 < z < 2.0$. |
Our results thus predict that the true increase in the proportion of high LEW systems with increasing recdshilt is significantly. greater than found by ? [rom the population of absorbers. | Our results thus predict that the true increase in the proportion of high EW systems with increasing redshift is significantly greater than found by \citet{2005ApJ...628..637N} from the population of absorbers. |
We have discussed the effects of dusty high. EW absorbers on the completeness of optical magnitucle-limited quasar samples over a range of absorber redshifts. | We have discussed the effects of dusty high EW absorbers on the completeness of optical magnitude-limited quasar samples over a range of absorber redshifts. |
The highest redshift. absorbers have been observed. using CRB optical alterelows as background. sources. (?).. | The highest redshift absorbers have been observed using GRB optical afterglows as background sources \citep{2009GCN..9215....1O}. |
Although GRBs would individually experience the same obscuration effects. CGlItDs are found to have a huge range in intrinsic brightness and spectra are obtained at variable time intervals following the GIs peak brightness. | Although GRBs would individually experience the same obscuration effects, GRBs are found to have a huge range in intrinsic brightness and spectra are obtained at variable time intervals following the GRB's peak brightness. |
The consequent very large dispersion in apparent GARB brightness. rather than the presence of modest amounts of dust in any intervening absorbers. dictates whether an object is observed. and absorption svstenis identified. | The consequent very large dispersion in apparent GRB brightness, rather than the presence of modest amounts of dust in any intervening absorbers, dictates whether an object is observed and absorption systems identified. |
Phorefore. one would expect a larger number of moderate {οV) absorption systems to be present in GIU spectra than quasar spectra. | Therefore, one would expect a larger number of moderate $E(B-V)$ absorption systems to be present in GRB spectra than quasar spectra. |
? identified 14 strong intervening ssystems (at a mean redshift of ἐς 1.1) along 14 CRB sight lines. an incidence roughly four times higher than along sight lines to quasars. | \citet{2006ApJ...648L..93P} identified 14 strong intervening systems (at a mean redshift of $\langle z \rangle=1.1$ ) along 14 GRB sight lines, an incidence roughly four times higher than along sight lines to quasars. |
The result is not expected if both GRBs and quasars sample random lines of sight. | The result is not expected if both GRBs and quasars sample random lines of sight. |
Since the intervening absorption svstems are thought. to be independent. of the background. source. the observed discrepancy. has led to a call to review the fundamental assumptions that underpin cxtragalactic absorption line research. | Since the intervening absorption systems are thought to be independent of the background source, the observed discrepancy has led to a call to review the fundamental assumptions that underpin extragalactic absorption line research. |
The discrepancy was confirmed. by 2.. but. the amplitude has since been reduced to a factor of 2.1+0.6 using a larger sample of 22 absorbers (?).. | The discrepancy was confirmed by \citet{2007ApJ...669..741S}, but the amplitude has since been reduced to a factor of $2.1 \pm 0.6$ using a larger sample of 22 absorbers \citep{2009A&A...503..771V}. |
A series of explanations have been suggested: to explain the observed CliD/quasar. discrepancy. | A series of explanations have been suggested to explain the observed GRB/quasar discrepancy. |
? claim that strong aabsorbing gas may be intrinsic to the CRB circumburst environment or originate from supernova remnants Wing in the same star-forming region. | \citet{2007ApJ...659..218P} claim that strong absorbing gas may be intrinsic to the GRB circumburst environment or originate from supernova remnants lying in the same star-forming region. |
It is also claimed (277) that source amplification due to strong eravitational lensing may bias the GRB spectral samples toward targets that contain more intervening absorbers. | It is also claimed \citep{2006ApJ...648L..93P,
2007ApJ...659..218P, 2009ApJ...706.1309T} that source amplification due to strong gravitational lensing may bias the GRB spectral samples toward targets that contain more intervening absorbers. |
Another mechanism has been proposed. (τὸν. which claims that the discrepancy. is due to the different bean. sizes of GRBs and quasars. but this has subsequently. been ruled out. by observational analysis (7).. | Another mechanism has been proposed \citep{2007Ap&SS.312..325F}, which claims that the discrepancy is due to the different beam sizes of GRBs and quasars, but this has subsequently been ruled out by observational analysis \citep{2007MNRAS.381L..99P}. |
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