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which is solved in the operator-splitiüng fashion. | which is solved in the operator-splitting fashion. |
Second. we have | Second, we have |
an ionisation parameter £230 erg em + at r—5 pe. a value at which OVIL and OVILE would both be prominent. | an ionisation parameter $\xi$$\sim$ 30 erg cm $^{-1}$ at r=5 pc, a value at which OVII and OVIII would both be prominent. |
In fact. the detection of Lle-like NVI and κο SiXIV in the RS spectrum shows the emitting eas must cover à wide range of tonisation parameter. those ions having peak abundances at £ of —5 erg em 1 and ~400 erg em . respectively. | In fact, the detection of He-like NVI and H-like SiXIV in the RGS spectrum shows the emitting gas must cover a wide range of ionisation parameter, those ions having peak abundances at $\xi$ of $\sim$ 5 erg cm $^{-1}$ and $\sim$ 400 erg cm $^{-1}$, respectively. |
Assuming the ionised Si and Ca emission lines detected in the MOS spectrum arise in the same outflow. requires still more highly ionised gas. while the excess Hux in the EPIC data at 6.6-6.SkeV. if correctly attributed to emission from FeXXV. extends that range to €~ L000 erg ems +. | Assuming the ionised Si and Ca emission lines detected in the MOS spectrum arise in the same outflow requires still more highly ionised gas, while the excess flux in the EPIC data at $\sim$ 6.6-6.8keV, if correctly attributed to emission from FeXXV, extends that range to $\xi$$\sim$ 1000 erg cm $^{-1}$. |
Such a broad ionisation structure. where the distance to the ionising source is constrained to rere.pe. suggests an inhomogeneous gas with a correspondingly broad. range of density. | Such a broad ionisation structure, where the distance to the ionising source is constrained to $\ga$ $_{BLR}$, suggests an inhomogeneous gas with a correspondingly broad range of density. |
While the present. spectra are clearly inadequate to allow such a complex ionised. outllow to be mapped it is possible to estimate several additional parameters. | While the present spectra are clearly inadequate to allow such a complex ionised outflow to be mapped it is possible to estimate several additional parameters. |
Line ratios in the Lle-like triplet are an establised diagnostic of ionised. plasmas. | Line ratios in the He-like triplet are an establised diagnostic of ionised plasmas. |
In the RS spectrum of the OVILE triplet is well defined. with the forbidden. (E). resonance (r) and intercombination (0) lines alb. detected (heure. 12). | In the RGS spectrum of the OVII triplet is well defined, with the forbidden (f), resonance (r) and intercombination (i) lines all detected (figure 12). |
The observed dine strengths are in the ratio 22:90:14 vielding values of the diagnostic parameters {Ξένα and CG(2fliír)-3. rather high and. low. respectively. compared with the calculated ratios for a pure. ow density photoionised plasma (Porquet and Dubau 2000). | The observed line strengths are in the ratio 22:9:4 yielding values of the diagnostic parameters $\sim$ 5 and $\sim$ 3, rather high and low, respectively, compared with the calculated ratios for a pure, low density photoionised plasma (Porquet and Dubau 2000). |
Alore recent calculations. which take account of the column density and ionisation parameter in the raciating eas (Cocet 22004). νιοα τὸ for OVILE for a. column. density of order μοι.E77 ? and £ = 30-100 erg enm | More recent calculations, which take account of the column density and ionisation parameter in the radiating gas (Godet 2004), yield G=3 for OVII for a column density of order $_{H}$$\sim$$10^{21-22}$ $^{-2}$ and $\xi$ = 30-100 erg cm $^{-1}$. |
Iincouraginglv. that column is of the same order as the product n.r in the above simple fit to the calculated emission measure. and also to the estimated. scattering z for a covering factor of 0.1. | Encouragingly, that column is of the same order as the product $_{e}$ r in the above simple fit to the calculated emission measure, and also to the estimated scattering $\tau$ for a covering factor of 0.1. |
In. particular the resonance line of OVILE (21.6 A)) will be optically thick in the core. limiting the elfeets of photo-excitation in comparison with higher order lines. | In particular the resonance line of OVII (21.6 ) will be optically thick in the core, limiting the effects of photo-excitation in comparison with higher order lines. |
We note also that the observed. Ér line ratios are lower in the Ne and Ale triplets. suggesting the opacity ellects there are less. | We note also that the observed f:r line ratios are lower in the Ne and Mg triplets, suggesting the opacity effects there are less. |
The mass of the (observed) extended gas envelope is ~900AL... assuming a uniform density. | The mass of the (observed) extended gas envelope is $\sim$, assuming a uniform density. |
With a projected oulllow velocity. observed at ~45 degrees to the line of sight. of ~700 kms +. the mass outflow rate is then. ~0.04 ALoveto with. an associated. mechanical. energy of "2.107"( erg +. These estimates would be reduced if the emitting gas is indeed. clunipy. | With a projected outflow velocity, observed at $\sim$ 45 degrees to the line of sight, of $\sim$ 700 km $^{-1}$, the mass outflow rate is then $\sim$ 0.04 $^{-1}$, with an associated mechanical energy of $2\times10^{40}$ erg $^{-1}$ These estimates would be reduced if the emitting gas is indeed clumpy. |
From the absorption-corrected 2.10 keV. Luminosity of 1.510 ere we estimate a bolometric luminosity [or oof 4«1075 erg | From the absorption-corrected 2–10 keV luminosity of $1.5\times 10^{43}$ erg $^{-1}$ we estimate a bolometric luminosity for of $4\times 10^{44}$ erg $^{-1}$. |
At an accretion clliclency of 0.1 the accretion. mass rate is. then ~0.06. 7. lo | At an accretion efficiency of 0.1 the accretion mass rate is then $\sim$ 0.06 $^{-1}$. |
pThus. as may e&encerallv be the case in Sevfert 1 galaxies. we find the ionised outflow in ccarrics a significant mass loss from the ACGN. | Thus, as may generally be the case in Seyfert 1 galaxies, we find the ionised outflow in carries a significant mass loss from the AGN. |
However. at least by a racial distance εκ τρ. the kinetic energy in the outllow is relatively small. even compared to the soft X-ray luminosity. | However, at least by a radial distance $\ga$ $_{BLR}$, the kinetic energy in the outflow is relatively small, even compared to the soft X-ray luminosity. |
EEPIC observations of high statistical qualit. confirm previous findings that the extremely hard (3.10 keV) power law index in the Sevfert 2 galaxy iis cue to strong continuum rellection from "cold matter. | EPIC observations of high statistical quality confirm previous findings that the extremely hard (2–10 keV) power law index in the Seyfert 2 galaxy is due to strong continuum reflection from `cold' matter. |
The intrinsic. power law. which is seen through a near-Compton-thick absorber. emerges above ~S keV ancl has a photon index tvpical of Sevfert 1 galaxies. | The intrinsic power law, which is seen through a near-Compton-thick absorber, emerges above $\sim$ 8 keV and has a photon index typical of Seyfert 1 galaxies. |
The dominant rellection component in aappears to be unallected by the large absorbing column. allowing [luorescent Line emission to be detected. from. Ni Ίνα (1.5 keV) down to Mg Wa (—1.25keV ). | The dominant reflection component in appears to be unaffected by the large absorbing column, allowing fluorescent line emission to be detected from Ni $\alpha$ $\sim$ 7.5 keV) down to Mg $\alpha$ $\sim$ 1.25keV). |
The combination of a [large column density obscuring 1e continuum source with the visibility of a large area of cold rellector suggests; that lis being viewed at an inclination which cuts the edge of 10 obscuring screen. | The combination of a large column density obscuring the continuum source with the visibility of a large area of cold reflector suggests that is being viewed at an inclination which cuts the edge of the obscuring screen. |
This would be consistent with the ecometry of “Sevlert galaxies on the edge. identified. in à recent extensive review of the optical polarisation properties of Sevlert galaxies (Smith 22002. 2004). | This would be consistent with the geometry of `Seyfert galaxies on the edge', identified in a recent extensive review of the optical polarisation properties of Seyfert galaxies (Smith 2002, 2004). |
Comparison with the sspectrum of NGC4051 (eg Schurch 22003) suggests that well-studied Sevlert 1: galaxy may be inherently similar but is being viewed from the other (lower obseuration) side of a (blurred) edge. | Comparison with the spectrum of NGC4051 (eg Schurch 2003) suggests that well-studied Seyfert 1 galaxy may be inherently similar but is being viewed from the other (lower obscuration) side of a (blurred) edge. |
The strong Fe Ίνα line appears to be cblue-shiftect in both EPIC pn and MOS spectra. indicating an origin in low ionisation matter with a projected outflow: velocity of 3100-1100 km t. | The strong Fe $\alpha$ line appears to be `blue-shifted' in both EPIC pn and MOS spectra, indicating an origin in low ionisation matter with a projected outflow velocity of $\pm$ 1100 km $^{-1}$. |
Although Less. well resolved: alter the inclusion. of a Compton shoulder. the Fe Ίνα line has a formal width of 320041600 km s (FWIIAL). | Although less well resolved after the inclusion of a Compton shoulder, the Fe $\alpha$ line has a formal width of $\pm$ 1600 km $^{-1}$ (FWHM). |
The projected. velocity and indicated velocity dispersion are both inconsistent with an origin at the far inner wall of a torus of rl pe. | The projected velocity and indicated velocity dispersion are both inconsistent with an origin at the far inner wall of a torus of $\ga$ 1 pc. |
Ho seems more likely that a large fraction of the neutral Wea emission arises in dense matter circulating or outllowing at a radial distance more typical of the BL clouds. | It seems more likely that a large fraction of the neutral $\alpha$ emission arises in dense matter circulating or outflowing at a radial distance more typical of the BLR clouds. |
For comparison. we note that Tran (1995) found the EWLIN of 113 in polarised light to be 6000 km inD | For comparison, we note that Tran (1995) found the FWHM of $\beta$ in polarised light to be 6000 km $^{-1}$ in. |
ES Below ~8 Κον a csoft excess emerges above (the hard. rellection-dominated. continuum. | Below $\sim$ 3 keV a `soft excess' emerges above the hard, reflection-dominated continuum. |
Spectral structure is resolved in both EPIC and HOS data. with the MOS spectrum showing remarkable detail. | Spectral structure is resolved in both EPIC and RGS data, with the MOS spectrum showing remarkable detail. |
In this soft X-ray band the spectrum is found to be dominated by resonance Lincs of He- and LH-like Si. Meg. Ne. O and N. with an observed bluc-shift of 470470 kms +. indicating an origin in a highly ionised. outllow. extending above the Sevfert 2. absorbing screen. | In this soft X-ray band the spectrum is found to be dominated by resonance lines of He- and H-like Si, Mg, Ne, O and N, with an observed blue-shift of $\pm$ 70 km $^{-1}$, indicating an origin in a highly ionised outflow extending above the Seyfert 2 absorbing screen. |
Relative line fluxes and the detection. of narrow (low temperature) radiative recombination continua of OVIL. OVILL ancl (probably) CVI are all consistent with the gas | Relative line fluxes and the detection of narrow (low temperature) radiative recombination continua of OVII, OVIII and (probably) CVI are all consistent with the gas |
16 LMC. | the LMC. |
This is the cohuun density projected perpendicular to the plane of 1ο LMC and is equivalent to the Ny—Nsin|b| measurcmeuts in the Calas. | This is the column density projected perpendicular to the plane of the LMC and is equivalent to the $N_\perp \equiv N \sin |b|$ measurements in the Galaxy. |
The first incasuremenuts of Galactic halo owards extragalactic objects found GV(Ovisin|bj=11.29 (8% staudard aeviation). | The first measurements of Galactic halo towards extragalactic objects found $\langle N(\mbox{\ovi})
\sin |b| \rangle = 14.29$ $38\%$ standard deviation). |
The average LAIC aud Ally Wav ΑΟvi) values and their y.andard deviations — are identical. | The average LMC and Milky Way $N_\perp (\mbox{\ovi})$ values – and their standard deviations – are identical. |
The observed kinematic profiles of the LMC aare quite broad. with breadths implying Tz(2/5)«109 K. Figure 2. shows the normalized aabsorptiou profiles along the 12 sight lines studied aud the profiles of a ioderate-streneth ttrausition. | The observed kinematic profiles of the LMC are quite broad, with breadths implying $T \lsim (2-5)\times10^6$ K. Figure \ref{fig:profiles} shows the normalized absorption profiles along the 12 sight lines studied and the profiles of a moderate-strength transition. |
The latter traces gas associated with neutral material iu the LMC disk. | The latter traces gas associated with neutral material in the LMC disk. |
The Hs ιο broader than the aalong all sight lines aud is svstematically bluc-shifted from the disk (bv ~30 oon average). | The is much broader than the along all sight lines and is systematically blue-shifted from the disk (by $\sim30$ on average). |
Thus. the ls kinematically decoupled from the bulk of the LAIC disk material. | Thus, the is kinematically decoupled from the bulk of the LMC disk material. |
There is eas present at velocities compatible with the outflow of material from the LMC disk alone all of the sight lines discussed by Howl ct al. | There is gas present at velocities compatible with the outflow of material from the LMC disk along all of the sight lines discussed by Howk et al. |
In only two | In only two |
where weivtan!Hnte;)/ReClin)] ls the ohase of the sigual from the sky aud Wyean![hnON;,)/ReCN,,)] is the phase of the noise. | where $\Psi_{lm}^{sig}=\tan^{-1}[{\bf Im}(c_{lm})/{\bf Re}(c_{lm})]$ is the phase of the signal from the sky and $\Psi_{lm}^{noise} =
\tan^{-1}[{\bf Im} (N_{lm})/{\bf Re}(N_{lm})]$ is the phase of the noise. |
Provided that CAIB anisotropy is a random Catssian field. he ay, of eq. | Provided that CMB anisotropy is a random Gaussian field, the $a_{lm}$ of eq. |
2 coefficient is a random variable with zero nean (iins= 0) aud variaLO nyyp=dyà"nanC1) where à, is the standard Kronecker svaubol aud C(7) the »ower spectu. | \ref{eq:eq2} coefficient is a random variable with zero mean $\langle a_{lm} \rangle=0$ ) and variance $\langle a_{lm} a_{l^{'}m^{'}} \rangle
=\delta_{ll^{'}}\delta_{mm^{'}} C(l)$ where $\delta_{lm}$ is the standard Kronecker symbol and $C(l)$ the power spectrum. |
Actually. he realization of the raudon CXMB signal on the sphere Is nique. which meaus that ina AT(0.o0) distribution ou he skv we have a siuele realization of the phases ouY. | Actually, the realization of the random CMB signal on the sphere is unique, which means that in a $\Delta T(\theta,\phi)$ distribution on the sky we have a single realization of the phases only. |
We denote the combines sieral by 5S5: aud the power spectrum of lis unique realization bv SS,|?. | We denote the combined signal by $S_{lm}$: and the power spectrum of this unique realization by $|S_{lm}|^2$. |
Qur task is to extract tje information about the beau slape either using Vy. Shin*lP or theirB combination. | Our task is to extract the information about the beam shape either using $\Psi_{lm}$ , $|S_{lm}|^2$ or their combination. |
.+ The manifestation of the beam shape asviuuetry in the Sin{?*D can be demoustrated inH the following+ way. | The manifestation of the beam shape asymmetry in the $|S_{lm}|^2$ can be demonstrated in the following way. |
Let us consider the following fiction Qualitatively the properties of tl| function A2) are the following. | Let us consider the following function Qualitatively the properties of the function $\Delta^2_s(lm)$ are the following. |
For simall / the CAD signals surpass the noise and the value of+ η* Dois letermined by the CXIB sieial. so for. this+ regionB (Asin/Ppo«I. | For small $l$ the CMB signals surpass the noise and the value of $|S_{lm}|^2$ is determined by the CMB signal, so for this region $\langle \Delta^2_s(lm) \rangle \ll 1$. |
For larger Si]? )»ecomes smaller. | For larger $l$, $|S_{lm}|^2$ becomes smaller. |
Because of t106 COSLLUC variance f.(fluctuatious of+ δρ]*P as a spectrum of the realization of a raxO11 process). SOIL Saul*2 σα1i become close to ZoYO. 5n,[P<Pεἰ where e is+ a constan aud e« |. | Because of the cosmic variance (fluctuations of $|S_{lm}|^2$ as a spectrum of the realization of a random process), some $|S_{lm}|^2$ can become close to zero, $|S_{lm}|^2<\epsilon^2$ where $\epsilon$ is a constant and $\epsilon\ll1$ . |
This meaus that Az(In) at these poiuts on tie (C;in j-plaue las Wanda The uunber deusitv of these maxima increases at larger f | This means that $\Delta^2_s(lm)$ at these points on the $(l,m)$ -plane has maxima The number density of these maxima increases at larger $l$. |
For ! of the order ἐν~o,!>1. where oO; FWIIM of the auteuua bea. the automa affects the spectrum. | For $l$ of the order $l_b\sim\Theta^{-1}_b\gg1$, where $\Theta_b\sim$ FWHM of the antenna beam, the antenna affects the spectrum. |
If the auteuua is asvuuuectric. this iuflueuce is also asviuietzic. | If the antenna is asymmetric, this influence is also asymmetric. |
Thus we have tle asviuinetric ΠΠion of anima in this region oft16 spectral planc. | Thus we have the asymmetric distribution of maxima in this region of the spectral plane. |
For verv largeo Ps. where the noise [Npin dondünates f nuniber density of maxima. Eq. (16)) | For very large $l$ 's, where the noise $N_{lm}$ dominates the number density of maxima, Eq. \ref{eq:Delta2}) ) |
is «etermuned by t ποσο, axd does not depend on the beam axd is therefore isotropic. | is determined by the noise, and does not depend on the beam and is therefore isotropic. |
The effects are demonstrated 1 Fig. 1.. | The effects are demonstrated in Fig. \ref{deltapower}, |
which i a result of a numerical experiment. | which is a result of a numerical experiment. |
Iu order to show the bean effect from A(Eky). we add ou the syiunetric part. | In order to show the beam effect from $\Delta_s(k_x,k_y)$, we add on the symmetric part. |
There are 25 contour levels between Ay(hy.ky)=0 and 5«107. | There are 25 contour levels between $\Delta_s(k_x,k_y)=0$ and $5
\times 10^{-3}$. |
The coordinates represeut the flat skv approxination of the general case at the limit foi291 (described iu Section 5). | The coordinates represent the flat sky approximation of the general case at the limit $l,m \gg 1$ (described in Section 5). |
luterval of modes that are scusitive to the beam asvunuuetry starts from /~7, and soes to imfuitv iu ideal conditions (no pixel noise. 6V7j 0). | Interval of modes that are sensitive to the beam asymmetry starts from $l\sim l_b$ and goes to infinity in ideal conditions (no pixel noise, $\langle N^2\rangle=0$ ). |
For the real situation the lait of this iuerval is finite aud determined by the pixel noise. fein)BrinDnhynn|2 | For the real situation the limit of this interval is finite and determined by the pixel noise, $|a_{lm}B_{lm}|^2\sim| N_{lm}|^2$. |
Toextract the information about the beam shape frou Fie. | To extract the information about the beam shape from Fig. |
l oue needs. for exame. to draw the averaged iso-density of the distribution of. maxima. ofDX AZ satisfied⋅⋅ to Eq. (16)). | \ref{deltapower} one needs, for example, to draw the averaged iso-density of the distribution of maxima of $\Delta^2_s$ satisfied to Eq. \ref{eq:Delta2}) ). |
One of the possilde inetlkyds of drawing this is described in the next section. | One of the possible methods of drawing this is described in the next section. |
m lis section we will show how to estimate the autenua bea Lsipe using the information οςtained in the phase cüstrilition of he signal iu the nap. using a single reaization of f1ο pphases of all {η7) mnodes. | In this section we will show how to estimate the antenna beam shape using the information contained in the phase distribution of the signal in the map, using a single realization of the phases of all $({lm})$ modes. |
After the descripion of this pise method it will be clear that. iu the sile case of liuear response oftje antenna. Cassia CB sienal auc noise. the plase method is equivalent to the power specrum method described above, | After the description of this phase method it will be clear that, in the simple case of linear response of the antenna, Gaussian CMB signal and noise, the phase method is equivalent to the power spectrum method described above. |
However. as we ll ionedin the introduction. iu he general case these luehocs eive differcit sets of information. | However, as we mentioned in the introduction, in the general case these methods give different sets of information. |
For he phase analvsis we will xvturb the phases by aclue controlled white noise iuto tje imap. | For the phase analysis we will “perturb” the phases by adding controlled white noise into the map. |
We consider au onsemble AL( Mx 1) where each element of the ensenibe consists of the same realization of theCAIB signal and pixel noise plus a raioni realization of a controlled. white noise |,m with⋅ the variauce⋅ oz>=const and raudoni plases for cach (21) inodo. | We consider an ensemble $M$( $M
\gg 1$ ), where each element of the ensemble consists of the same realization of theCMB signal and pixel noise plus a random realization of a controlled white noise $W_p$ with the variance $\sigma_W^2=const$ and random phases for each $(lm)$ mode, |
estimate the power spectra. | estimate the power spectra. |
The dotted line in the figure represents w/k—5kms!. | The dotted line in the figure represents $\omega / k=5~\mathrm{km~s}^{-1}$. |
Since we use the subsonic filtering method, spectral power above the dotted line is significantly reduced. | Since we use the subsonic filtering method, spectral power above the dotted line is significantly reduced. |
The region in the lower frequency and the smaller wave number has greater power. | The region in the lower frequency and the smaller wave number has greater power. |
Integrating k-w diagram over wave number space, we can estimate the frequency power spectrum density, P,. | Integrating $k$ $\omega$ diagram over wave number space, we can estimate the frequency power spectrum density, $P_{\nu}$. |
The power spectrum density is defined as, where V, is the LCT velocity, νι is the lowest frequency determined by total duration, and 14,,, is the highest frequency coming from the observational sampling time. | The power spectrum density is defined as, where $V_{\perp}$ is the LCT velocity, $\nu_{min}$ is the lowest frequency determined by total duration, and $\nu_{max}$ is the highest frequency coming from the observational sampling time. |
The symbol, <>, denotes the temporal average. | The symbol, $<>$, denotes the temporal average. |
Figure 3 shows an example of the power spectrum density. | Figure \ref{tspectrum} shows an example of the power spectrum density. |
Power spectra can be fitted to double power-law function, P,οςv?* when v«νι (so-called break frequency) and P,ev" when v>1%. | Power spectra can be fitted to double power-law function, $P_\nu \propto \nu^{\alpha_L}$ when $\nu < \nu_b$ (so-called break frequency) and $P_\nu \propto \nu^{\alpha_H}$ when $\nu \ge \nu_b$. |
Typically, /«Vi>,%,or, and og of LCT velocity are 1.1 km s!, 4.7 mHz, -0.6, and -2.4, respectively in this work. | Typically, $\sqrt{<V_{\perp}>},\nu_b, \alpha_L, $ and $ \alpha_H$ of LCT velocity are 1.1 km $^{-1}$, 4.7 mHz, -0.6, and -2.4, respectively in this work. |
By using the same analysis, temporal power spectra of divergence and rotation of velocity fields are also estimated (the dashed line and the dotted line in figure 3)). | By using the same analysis, temporal power spectra of divergence and rotation of velocity fields are also estimated (the dashed line and the dotted line in figure \ref{tspectrum}) ). |
The total power, ar, and ag of the divergence is 8.4 x 10? s, -0.2, -1.5 respectively and the divergence will contribute to produce the acoustic mode waves. | The total power, $\alpha_L$, and $\alpha_H$ of the divergence is 8.4 $\times$ $^{-3}$ $^{-1}$, -0.2, -1.5 respectively and the divergence will contribute to produce the acoustic mode waves. |
The total power, ay, and og of the rotation is 6.8 x 10? s1, -0.2, -1.3 respectively and the rotation is related to the generation of torsional Alfvénn waves. | The total power, $\alpha_L$, and $\alpha_H$ of the rotation is 6.8 $\times$ $^{-3}$ $^{-1}$, -0.2, -1.3 respectively and the rotation is related to the generation of torsional Alfvénn waves. |
The power spectra of the LCT velocity has a harder power law index than that of its derivatives (V-V,Vx V). | The power spectra of the LCT velocity has a harder power law index than that of its derivatives $\nabla \cdot V, \nabla \times V$ ). |
'The resulting power spectra are quite essential for evaluating models of coronal heating by Alfvénn waves (Hollweg1982;Kudoh&Shibata1999;SuzukiInutsuka2005,2006;Matsumoto&Shibata 2010). | The resulting power spectra are quite essential for evaluating models of coronal heating by Alfvénn waves \citep{holl82,kudo99,suzu05,suzu06,mats10}. |
. In their numerical simulations, nonlinear Alfvénn waves are driven by photospheric convection to heat the corona. | In their numerical simulations, nonlinear Alfvénn waves are driven by photospheric convection to heat the corona. |
Matsumoto&Shibata(2010) include the observed temporal power spectra shown here to generate Alfvénn waves and succeed in producing sufficient energy flux to heat the corona. | \cite{mats10} include the observed temporal power spectra shown here to generate Alfvénn waves and succeed in producing sufficient energy flux to heat the corona. |
The combination band 153) is surprisingly strong in the laboratory spectrum of Tulejetal.(1998).. and it was suggested that this band may correspond to the 44964 DIB. | The combination band $1^2_03^1_0$ is surprisingly strong in the laboratory spectrum of \citet{tulej}, and it was suggested that this band may correspond to the $\lambda$ 4964 DIB. |
Since the 153], band was not revisited in the experiment of Lakinetal.(2000).. we cannot examine in detail its agreement with the A4964 DIB. | Since the $1^2_03^1_0$ band was not revisited in the experiment of \citet{lakin}, we cannot examine in detail its agreement with the $\lambda$ 4964 DIB. |
However. with our substantially larger sample of stars. we are in à position to re-examine the correlation between the intensities of A4964 and A6270 (supposedly the origin band of C5). | However, with our substantially larger sample of stars, we are in a position to re-examine the correlation between the intensities of $\lambda$ 4964 and $\lambda$ 6270 (supposedly the origin band of ). |
If these two bands are due to the same species. they must have the same intensity ratio from star to star. as this ratio is determined solely by the Franck-Condon factors. | If these two bands are due to the same species, they must have the same intensity ratio from star to star, as this ratio is determined solely by the Franck-Condon factors. |
Figure 6. displays the spectra of A4964 and A6270 in a sample of twelve reddened stars. | Figure \ref{band120310} displays the spectra of $\lambda$ 4964 and $\lambda$ 6270 in a sample of twelve reddened stars. |
While it appeared in our original work (MeCall.York.andOka2000) that these bands were correlated. this was apparently due to the small sample (4) of stars considered in that work. | While it appeared in our original work \citep{myo} that these bands were correlated, this was apparently due to the small sample (4) of stars considered in that work. |
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